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1403 | 1403.0822.txt | Stars are born in dense cores of molecular clouds. The core mass function (CMF), which is the mass distribution of dense cores, is important for understanding the stellar initial mass function (IMF). We obtained 350 \micron\ dust continuum data using the SHARC-II camera at the Caltech Submillimeter Observatory (CSO) telescope. A 350 \micron\ map covering 0.25 ${deg}^{2}$ of the Ophiuchus molecular cloud was created by mosaicing 56 separate scans. The CSO telescope had an angular resolution of 9 \arcsec, corresponding to $1.2\times {10}^{3}\ $AU at the distance of the Ophiuchus molecular cloud (131 pc). The data was reduced using the Comprehensive Reduction Utility for SHARC-II (CRUSH). The flux density map was analyzed using the GaussClumps algorithm, within which 75 cores has been identified. We used the Spitzer c2d catalogs to separate the cores into 63 starless cores and 12 protostellar cores. By locating Jeans instabilities, 55 prestellar cores (a subcategory of starless cores) were also identified. The excitation temperatures, which were derived from FCRAO ${}^{12}$CO data, help to improve the accuracy of the masses of the cores. We adopted a Monte Carlo approach to analyze the CMF with two types of functional forms; power law and log-normal. The whole and prestellar CMF are both well fitted by a log-normal distribution, with $\mu =-1.18\pm0.10,\ \sigma =0.58\pm0.05$ and $\mu =1.40\pm0.10,\ \sigma =0.50\pm0.05$ respectively. This finding suggests that turbulence influences the evolution of the Ophiuchus molecular cloud.% Conversely, the starless CMF is better fitted by a power law distribution, with $\alpha =-1.51\pm 0.20$. \\%The measured CMFs can be well fitted by log-normal functions. This finding suggests that turbulence influences the evolution of the Ophiuchus molecular clouds. % 46 cores' mean temperatures were estimated by combined our 350 \micron\ data with SCUBA/JCMT 850 \micron\ data. %The starless CMF is similar to the IMF of Salpeter (1955), which means that maybe the mass of stars is set at the starless stage in the core formation process. | \no Molecular cores are dense condensations within molecular clouds, in which stars are born [1-3]. The study of the CMF is therefore potentially important for understanding the$\ $IMF. The similarity between the CMF and the IMF suggests that$\ $stars may form at the same efficiency in all cores\ [4], although this remains unclear [5,6]. Protostellar cores have embedded self-luminous sources. Protostellar cores have lost some mass because of accretion onto the embedded protostar, or the extistence of bipolar outflows [7-9]. These changes could mean that protostellar cores no longer exhibit the same initial conditions as starless cores. Therefore, in %\altaffiltext{1}{Accepted by SCIENCE CHINA Physics, Mechanics \& Astronomy.} %\altaffiltext{2}{E-mail contact: [email protected]} \vspace*{-3.0mm} \noindent\rule{2.5cm}{0.4pt}\\[0.1mm]{\qihao *Corresponding author (email: [email protected])\\Accepted by SCIENCE CHINA Physics, Mechanics \& Astronomy.}%ÊÖ¶¯E-mailµØÖ· %\noindent\rule{2.5cm}{0.4pt}\\[0.1mm]{\qihao Accepted by SCIENCE CHINA Physics, Mechanics \& Astronomy.}%ÊÖ¶¯E-mailµØÖ· %\vspace*{1mm} %\noindent\rule{2.5cm}{0.4pt}\\[0.1mm] %{\footnotesize{ %\hspace*{3.5mm}*Corresponding author (email: [email protected]) %}} \no order to better characterize the initial conditions of star formation, the differentiation of protostellar cores from starless cores needs to be investigated [10]. Herein we use both power law and log-normal functional form to fit the CMF. The power law form is the traditional form of the IMF, but in recent years it has been proposed that the low mass region of the CMF may be characterized by a log-normal function [10]. Ballesteros-Paredes et al. [11] calculated the CMF using numerical models of turbulent fragmentation of molecular clouds, resulting in a form that does not follow a single power-law, and instead is more similar to a log-normal function. Padoan and Nordlund [12] have deemed that the slope of the CMF depends on the slope of the turbulent energy spectrum. Previous research has shown that a log-normal distribution arises when the central limit theorem is applied to isothermal turbulence [13-15]. The aim of this study is to determine the CMF of the Ophiuchus molecular cloud, so as to advance the current understanding of how cores evolve into stars in a medium mass star forming region such as the Ophiuchus molecular cloud. We used core extraction and the CMF fitting techniques similar to those in Li et al. [16], which studied high mass quiescent cores in the Orion molecular cloud. The Ophiuchus molecular cloud is a medium mass star forming region [17]. There is no evidence for the existence of massive star clusters, H II regions and other usual marks of massive star formation, such as water masers, but there is one B star present within the cloud. There are several previously published dust continuum maps of the Ophiuchus molecular cloud. Motte et al. [18] mapped approximately 0.13 ${deg}^{2}$ at 1.3 mm, with an effective angular resolution of 15 \arcsec. The prestellar CMF can be fitted by a power law, $\rm dN/dM\sim {M}^{-1.5}$ for $\rm M <0.5$ \Ms, while the whole CMF appears to steepen to $\rm dN/dM\sim {M}^{-2.5}$ for $\rm M >0.5$ \Ms. Johnstone et al. [19] mapped approximately 0.19 \squaredeg\ at 850 \micron, with an angular resolution of 15 \arcsec, identifying 55 cores. The CMF is fitted by two power laws, $\rm N(M)\sim {M}^{-\alpha }$, with $\alpha =1.0-1.5$ for $\rm M >0.6$ \Ms, and $\alpha =0.5$ for $\rm M \leq 0.6$ \Ms. Young et al. [20] mapped approximately 11 \squaredeg\ at 1.1 $mm$, with an angular resolution of 31 \arcsec, identifying 44 cores and a few candidates for cores. The CMF can be fitted by a power law function, $\rm N(M)\sim {M}^{-\alpha }$, with $\alpha =2.1\pm 0.3$ for $\rm M >0.5$ \Ms . Stanke et al. [21] mapped approximately 1.3 \squaredeg\ at 1.2 $mm$, with an angular resolution of 24 \arcsec, identifying 143 cores in this map, including 111 starless cores. The starless CMF is fitted by a broken power law, where the power law indices are -0.3, -1.3, and -2.3. Herein we use the data of dust continuum at 350 \micron, which is close to the peak of the dust spectral energy distribution, to map approximately 0.25 \squaredeg\ at this wavelengh. The CSO telescope provides 9 \arcsec angular resolution at 350 \micron\ which is 3 times better than Herschel [22]. Johnstone et al. [18] found 55 cores within a SCUBA/JCMT 850 \micron\ map of the Ophiuchus cloud. 54 cores reside within our 350 \micron\ map. They are indicated with the blue square symbol in Fig. 10. The cores unique to the 350 \micron\ map may have too high a temperature to have been detected by the Johnstone 850 \micron\ survey. Some of the cores apparently missing from this 350 \micron\ study in comparison with the 850 \micron\ study are likely to be cooler sources. % %The cores unique to the 350 \micron\ map maybe have too high a temperature to be detected by the 850 \micron\ survey. Some of the cores apparently missing from this 350 \micron\ study are likely to be cooler sources. %The cores?with?stars?should?deviate?significantly?as?some?of?the?mass?has?been?lost?to?the?star?and?possibly?an?outflow. In order to study the initial conditions of star formation better, we first must divide the cores into: starless cores, protostellar cores and presteller cores. Protostellar cores can be found by looking at the Spitzer c2d catalogues. Prestellar cores can be distinguished by their property of Jeans instability. We made a detailed introduction in part 3.5. %Since dense cores are thought to be the precursors of proto-stars (e.g., Enoch et al. 2006) and the core mass function (CMF) of such cores is similar in shape to the stellar initial mass function (Motte et al.1998\ [28], 2001\ [29]; Alves et al. 2007;K¡§ onyves et al.2010), %The starless CMF can be well fitted by the power law forms. But the exponents ($\alpha $) of those power law form CMF are significantly greater than exponents the Salpeter IMF($\alpha =-2.35$). Clark et al. (2007)[38] raised that the CMF depends on the cores lifetime, which makes the exponent of the power law form IMF significantly less than that of the CMF. The starless CMF is related to the process of star formation and the IMF. The process of star formation determines the speed of each starless cores evolution at a certain mass. The IMF determines the relative number of stars after the star formation is complete [35]. The starless CMF has a shape similar to the stellar IMF. It is likely to assume that the two distributions are related, e.g. by the star formation efficiency, which quantifies how cores are converted into stars. At low mass end, sensitivity limits cause a reduction in the number of cores found. This Completeness more or less affects the shape of the CMF. The CMF is found to be similar to the form of the stellar initial mass function (IMF). Maybe mass of stars is set at the starless stage during the core formation process (e.g. Padoan $\&$ Nordlund 2002 [41], [60], [1]). Such a connection between the CMF and IMF would appear to be a natural prediction of theories based on the inside-out collapse of gravitationally bound cores (Shu 1977 [42]; McKee \& Tan 2002[43]), but inconsistent with the theory of competitive accretion(e.g.,Bonnell et al. 1997[44]). This suggests that the form of the IMF is the result of the form of the starless CMF. The starless CMF provides the origin of the IMF. % %Turbulence within molecular clouds is highly supersonic, which plays a major role in the process of star formation (see e.g Ballesteros-Paredes et al. 2007 [47], McKee \&Ostriker 2007 [49], Joos et al. 2013[48]). In molecular clouds, large-scale supersonic turbulence cascades to small scales by shocks [36]. The probability density function (PDF) can describe the gas density fluctuations created by turbulence. A log-normal probability distribution in density was naturally predicted, when the central limit theorem is applied to isothermal turbulence (Larson 1973 [7]; Adams\& Fatuzzo 1996 [8]). And this has been produced in computer simulations (e.g., Vazquez-Semadeni 1994[41]; Klessen 2001[50]). Padoan \& Nordlund 2002 [5] first derived the mass distribution of gravitationally unstable cores with a log-normal PDF of of the mass density in supersonic turbulence, and concluded that turbulent fragmentation is essential to the origin of the stellar IMF. Such a log-normal behaviour for density fluctuations has received observational support. %The jeans' length is about $1.5\times {10}^{3}$ Au. %The thermal Jeans scale is about 0.44 Au. %\vspace*{1mm} %\noindent\rule{2.5cm}{0.4pt}\\[0.1mm]{\qihao *Corresponding author (email:[email protected])}%ÊÖ¶¯E-mailµØÖ· %The cold temperatures and high density allow gravity to overcome thermal pressure, initiating the collapse toward new stars %The observational similarity between the CMF and the IMF was rst put forth by Motte et al. (1998), and since this time many other samples of dense cores have been presented in this context (e.g., Simpson et al. 2008; Enoch et al. 2008; Alveset al. 2007; Nutter \& Ward-Thompson 2007; Stanke et al.2006; Johnstone et al. 2001). The qualitative similarity between the CMF and the IMF offers support for the accepted idea that stars form from dense cores (e.g. Ladaet al. 2008) | Using the SHARC-II camera at the CSO telescope, we obtained dust continuum data at 350 \micron\ for the Ophiuchus molecular cloud. These data were mainly used to determine the CMF of the Ophiuchus molecular cloud. The total masses, mean temperatures, Jeans Length, and mean mass density of the cores were also calculated.\\ %Using the SHARC-II camera at the CSO telescope, we obtained dust continuum data at 350 \micron\ of the Ophiuchus molecular cloud. These data were mainly used to determine the CMF of the Ophiuchus molecular cloud. Core's total mass, mean temperature, Jeans Length, mean mass density were also calculated. \\ We summarize the primary results of the paper as follows:\\ Firstly, A 350 \micron\ map covering 0.25 deg${}^{2}$ of the Ophiuchus molecular cloud was created by mosaicing 56 separate scans.\\ Secondly, This map of the Ophiuchus molecular cloud was analyzed using the GaussClumps algorithm, in which 75 cores have been detected. The mean temperature of the cores was 29K and the mean mass of the cores was 0.6 M$_\odot$. \\ Thirdly, We used the Spitzer c2d catalogs to separate 63 starless cores from 12 protostellar cores. 55 prestellar cores were also distinguished by their property of Jeans instability.\\ Fourthly, We performed a comparison between the core catalogue at 350 \micron\ and the JCMT/SCUBA 850 \micron\ catalogue. The cores unique to the 350 \micron\ map may have too high a temperature to be detected by the 850 \micron\ survey. Some of the cores apparently missing from this 350 \micron\ study are likely to be cooler sources.\\ %Comparison between the core catalogue at 350 \micron\ against the JCMT/SCUBA 850 \micron\ catalogue. The cores unique to the 350 \micron\ map maybe have too high a temperature to be detected by the 850 \micron\ survey. Some of the cores apparently missing from this 350 \micron\ study are likely to be cooler sources.\\ Lastly, We found that the whole and prestellar CMF can be well fitted by log-normal distributions with $\mu =-1.18\pm0.10,\ \sigma =0.58\pm0.05$ and $\mu =1.40\pm0.10,\ \sigma =0.50\pm0.05$ respectively, while the starless CMF can be better fitted by a power law: $\rm dN/dM \sim {M}^{-1.51\pm 0.20},\ {M}_{max}=3.1\ M_\odot,\ {M}_{min}=0.11\pm 0.01\ M_\odot$.%, while the whole and prestellar CMF can be better fitted by log-normal.\\ %Clark et al.[41] raised that the CMF depends on the cores lifetime, which maybe makes the exponent ($-1.51\pm 0.20$) of the starless CMF is greater than the the exponent (-2.35) of the Salpeter IMF [36].\\ %Thanks for our paper's editor Guo Yuanyuan's warm, considerate and attentive editing service. After we adopted a lot of anonymous reviewer's suggestions, the quality and scientific importance of the paper significantly increased. Many thanks are due to the anonymous reviewer. \vspace*{2mm} | 14 | 3 | 1403.0822 |
1403 | 1403.3435_arXiv.txt | We present a sensitive search for \WS\ $W3$ ($12\micron$) and $W4$ ($22\micron$) excesses from warm optically thin dust around \hip\ main sequence stars within 75~pc from the Sun. We use contemporaneously measured photometry from \WS, remove sources of contamination, and derive and apply corrections to saturated fluxes to attain optimal sensitivity to $>10\micron$ excesses. We use data from the \WS\ All-Sky Survey Catalog rather than the AllWISE release, because we find that its saturated photometry is better behaved, allowing us to detect small excesses even around saturated stars in \WS. Our new discoveries increase by 45\% the number of stars with warm dusty excesses and expand the number of known debris disks (with excess at any wavelength) within 75~pc by 29\%. We identify 220 \hip\ debris disk-host stars, 108 of which are new detections at any wavelength. We present the first measurement of a 12$\micron$ and/or 22$\micron$ excess for 10 stars with previously known cold (50--100~ K) disks. We also find five new stars with small but significant $W3$ excesses, adding to the small population of known exozodi, and we detect evidence for a $W2$ excess around HIP96562 (F2V), indicative of tenuous hot (780~K) dust. As a result of our \WS\ study, the number of debris disks with known 10--30$\micron$ excesses within 75~pc (379) has now surpassed the number of disks with known $>30\micron$ excesses (289, with 171 in common), even if the latter have been found to have a higher occurrence rate in unbiased samples. | Numerous surveys have been conducted to search for dusty disks around main sequence stars over the last three decades. The all-sky survey performed by the \textit{Infrared Astronomical Satellite (IRAS)} was the first to detect infrared (IR) excess emission from circumstellar dust disks at 25 and 60$\mu m$, with $\sim$~170 disks identified in all. Subsequent pointed surveys with the \textit{Infrared Space Observatory} (\iso\ ), the \textit{Spitzer Space Telescope}, and the \textit{Herschel Space Observatory}, and the recent all-sky survey by the \akari\ satellite have greatly increased the number of disks discovered. To date, over 350 debris disks are known around main sequence stars within 75~pc \citep[e.g.,][and references therein]{Su2006, Moor2006, Moor2009, Moor2011, Bryden2006, Rhee2007, Trilling2008, Hillenbrand2008, Carpenter2009, Mizusawa2012, Fujiwara2013, Eiroa2013, Wu2013, Cruz-SaenzdeMiera2013}, and several hundred more around more distant stars, including open cluster members, out to $\sim$1~kpc (e.g., Siegler 2007, Currie 2008a,b)% Most ($\sim$85\%) of the known debris disks in the solar neighborhood are comprised of cold ($<$100~K) circumstellar dust. These have been identified through their characteristically strong emission at wavelengths longer than 30~$\micron$, at which the disks are often orders of magnitude brighter than the stellar photosphere. This cold dust is analogous to debris produced from destructive collisions in the solar system Edgeworth-Kuiper Belt (EKB). The dust has to be continually produced in such collisions because its lifetime in the system is short: large grains spiral into the star due to Poynting-Robertson drag, and small grains are blown outward by radiation pressure. Both processes remove dust on characteristic time scale shorter than one million years \citep{Backman1993}: much less than the ages of stars in the solar neighborhood. Except in cases of stars with obvious signatures of youth, the detection of cold circumstellar dust demonstrates the presence of a belt of colliding planetesimals which, like the dust, are likely located in the cold outer reaches of the system (i.e., $>$ 10~AU from the star). Most known, faint warm debris disks have been discovered from pointed surveys with \spitzer\ \citep[e.g.,][]{Su2006, Trilling2008, Carpenter2009}. Deep targeted observations with the \spitzer\ Infrared Spectrograph \citep[IRS;][]{Houck2004}, in particular, have allowed the measurement of excesses peaking in the 10--30~$\micron$ range at only 3\% of the photospheric flux at the same wavelengths \citep{Carpenter2009, Lawler2009}. The advantage in using 5--30~$\micron$ mid-IR spectroscopy is that it allows an accurate calibration of the stellar photospheric flux---essential for detecting small excesses. However, pointed surveys by design are limited in scope, and the data interpretation is subject to biases in the sample selection. \WS\ offers an opportunity to search for warm debris disks over the entire sky in an unbiased fashion. Though not as sensitive as deep, pointed Spitzer observations, \WS\ is 100--600 times more sensitive than {\it IRAS} and 10--50 times more sensitive than \akari\ in the mid-IR --- making it by far the most sensitive all-sky survey at these wavelengths. Through near-simultaneous and uniform 3--30~$\micron$ photometry, \WS\ also enables accurate calibration of the stellar photospheres, and hence good sensitivity to faint mid-IR excesses with $<10$\% of the 10--30~$\micron$ photospheric flux. Numerous searches of the \WS\ catalog have already been conducted to identify debris disks. \citet{Krivov2011}, \citet{Morales2012}, \citet{Ribas2012}, \citet{Lawler2012}, and \citet{Kennedy2012} sought $W3$ and $W4$ excesses among known extrasolar planet hosts. Approximately two dozen distinct planet-host stars with possible $W3$ or $W4$ excesses are found among these studies. \citet{Rizzuto2012}, \citet{Riaz2012}, \citet{Luhman2012}, and \citet{Dawson2013} sought \WS\ excesses in the young Scorpius-Centaurus association. The total number of disks identified in these studies is $\approx$160, with some duplications and/or non-confirmations among the three teams (note that not all of these were debris disks). Finally, \citet{Avenhaus2012}, \citet{Kennedy2013}, \citet{Wu2013}, \citet{Cruz-SaenzdeMiera2013}, and \citet{Vican2014} sought debris disks among solar neighborhood stars. \citet{Avenhaus2012} find no new $W3$ or $W4$ excesses around the 100 nearest M dwarfs. \citet{Kennedy2013} identify 15 known and 7 new $W3$ excesses around \hip\ stars within 150~pc. An excess at such relatively short wavelengths may indicate the presence of an exozodi: a dust population at a similar temperature to the solar system's zodiacal dust. The recent studies of \citet{Wu2013} and \citet{Cruz-SaenzdeMiera2013} are most similar to ours in design. \citet{Wu2013} seek $W4$ excesses around \hip\ stars of all spectral types within 200 pc, while \citet{Cruz-SaenzdeMiera2013} seek $W4$ excesses around F2--K0 stars brighter than $V=15$~mag. As we discuss in \S \ref{sec:comparison2Wu2013} and \ref{sec:comparison2Cruz2013}, our results are mostly complementary to the results from these studies. Importantly, through a careful calibration of \WS\ photometric systematics, we are able to detect excesses that are fainter than those reported in \citet{Wu2013} and \citet{Cruz-SaenzdeMiera2013}. Our newly-identified disk-host stars are also often either brighter (saturated in \WS) than those considered in \citet{Wu2013} and \citet{Cruz-SaenzdeMiera2013}, or fainter (with $W4$ SNR less than 20) than those considered in \citet{Wu2013}. An accurate understanding of \WS\ photometry systematics is essential to reliable identification of dust excesses. The strongest systematic effect is the over-estimation of the $W2$ fluxes of bright ($W2<6.7$~mag) stars from profile-fit photometry \citep[see \S~VI.3.c.i.4.\ of][]{Cutri2012}, but \citet{Kennedy2013} and several additional studies also note remnant offsets in the \WS\ photometry and colors that render some previously-reported tenuous excesses uncertain. We address this and other more subtle flux-dependent trends in the \WS\ photometry in \S~\ref{sec:wise_systematics}. Other reasons for mis-identifications include confusion with background IR-bright sources seen in projected proximity, contamination from interstellar cirrus, and unknown amounts of interstellar extinction. Various approaches have been adopted to mitigate these effects, including source position comparisons between the short- and long-wavelength \WS\ filters, exclusion of extended IR-bright regions in \iras, confirmation of excesses through spectral energy distribution (SED) fitting, and, importantly, visual inspection of the stellar images (e.g., Kennedy \& Wyatt 2013). We have incorporated all of these techniques, and others, in our approach (\S\ref{sec:sample_definition}), and furthermore have only selected candidates at confidence levels greater than 99.5\% or 98\% at $W4$ or $W3$ respectively, based on the \textit{empirical} scatter in WISE photometry. Importantly, we identify debris disk candidates using only \WS\ colors: the fact that these are homogeneous and simultaneous set of measurements reduces our vulnerability to stellar variability and other sources of error. Our results therefore present an opportunity for an unbiased analysis of the occurrence and evolution of warm circumstellar dusty disks. We describe the method we used to identify IR excesses in \S\ref{sec:iridentify}. We present our cross-match with the entire \hip\ catalog \citep{Perryman1997, vanLeeuwen2007} with the \WS\ All-Sky Catalog \citep[ASC;][]{Wright2010} and define our working sample of stars in \S\ref{sec:WHXD} and \S\ref{sec:sample_definition}, respectively. \S\ref{sec:atlas_vs_single} addresses a previously unknown issue that we discovered with the reliability of \WS\ ASC photometry on certain stars. In \S\ref{sec:wise_systematics}, we outline how we precisely calibrated the \WS\ photometric systematics to produce a set of reliable debris disk detections for stars in our sample. Section \ref{sec:IRAnalysis} describes our IR excess identification procedure. Section 2.6 describes our test with identifying IR excesses in the more recent, AllWISE data release, and presents our arguments for the higher reliability of bright-star photometry in the preceding All-Sky data release. In \S\ref{sec:modeling}, we describe our procedure for quantifying basic disk characteristics. Section \ref{sec:analysis} offers an analysis of the inferred circumstellar locations of the detected excesses: whether they belong to exozodi, asteroid belt analogs, or previously known colder EKB analogs. Section \ref{sec:discussion} discusses our results in the context of previous surveys with {\it IRAS}, {\it Spitzer}, {\it AKARI}, and \WS. | \label{sec:conclusion} We identify a volume-limited sample of \hip\ stars within 75 pc that show infrared excess fluxes based on photometry contained in the \WS\ All-Sky Data Release. We carefully screen the \WS\ photometry for various sources of false-positives both astrophysical and instrumental. One such issue, newly identified in our work, is that in a tiny fraction of \WS\ photometry, the median of single-exposure fluxes is inconsistent with the \WS\ All Sky Catalog flux, and neither is reliable. We reject photometry compromised by this and other issues; precisely calibrate flux-dependent systematic effects in saturated photometry; and correct for the dependence of \WS\ colors on photospheric temperature. Using the blue wing of the resulting color distributions to empirically evaluate our false positive rate (FPR) for the red outliers that correspond to dusty circumstellar disks, we robustly detect 215 such disks at 22$\micron$ with FPR $<$ 0.5\% and 5 additional disks at 12$\micron$ with FPR $<$ 2\%. Our careful screening and precise calibration of the \WS\ photometry enables us to identify faint circumstellar dust disks that had gone unnoticed in previous analyses, in addition to confirming disks that had been previously detected using photometry from \WS\ and other missions. Our new detections represent, in total, an increase of 45\% in the number of stars within 75~pc known to have flux excesses at mid-IR wavelengths. In contrast to {\it IRAS} and {\it ISO}, which produced many detections of cold circumstellar dust, the \WS\ mid-infrared bands have enhanced sensitivity to warmer dust in regions analogous to our own Solar System's asteroid belt and zodiacal cloud --- regions most likely responsible for terrestrial planet formation. We report the following detections: \begin{enumerate} \item 220 stars with FPR$<$0.5\% mid-IR excesses at 22$\micron$ and/or FPR$<$2\% excesses at 12$\micron$. For 113 of these we present the first detection of a debris disk at any wavelength, and for a futher 10 that have known longer-wavelength excesses, we present the first measurement of an excess at $12\micron$ and/or $22\micron$. \item A subset of 211 of our disks are detected with significant excesses in 22$\micron$ only. Aggregate 12$\micron$ excesses can be detected by weighted averages of the 12/22$\micron$ excess flux ratio over different subsets of this sample, and these aggregate 12$\micron$ detections are highly significant. The subset with previously published low (50--120 K) dust temperatures has an aggregate 12/22$\micron$ excess flux ratio consistent with low-temperature dust, while the aggregate flux ratio for the previously unknown disks indicates that many of them have dust at asteroidal temperatures ($>$ 130 K). \item A subset of 4 stars possess significant excess detections at both 12 and 22$\micron$, with a flux ratio indicative of dust temperatures ranging from $\sim 130$~K to $\sim 200$~K. All of these systems are known to possess long-wavelength ($>60\mu m$) excesses well fit by colder dust, and none were suspected to have $>$100~K dust. Hence, our results indicate the presence of dust at multiple temperatures in these systems. \item A subset of 5 disks are detected with signficant excesses only at 12$\micron$. Upper limits to the 22$\micron$ excesses in these systems yield 3$\sigma$ lower limits on the temperature ranging from $\sim 175$~K to $\sim 275$~K. While the coolest of these limits would permit asteroidal-temperature dust, the data are more consistent with warmer dust. Such dust would overlap with the habitable zones in these systems and could come from planetesimals left over from the formation of terrestrial planets. \item Five additional stars, not included in our count of 220 detected dust disks, possess shorter-wavelength excesses at 4.6$\micron$ with FPR$<$ 5\%. One of these excesses, around the F2V star HIP~96562, is suspected to be caused by hot (780~K) dust. The origin of the remaining four excesses, all associated with K dwarfs, remains speculative. It is possible that in two of the cases the thermal emission is caused by tenuous amounts of hot, short-lived, sub-micron-sized dust. However, this scenario can not account for all four cases of $W2$ emission from K stars. We therefore suggest an alternate explanation involving chromospheric activity. \item $1.8$\%$ \pm 0.2$\% of solar type (FGK) stars and $21.6$\%$ \pm 2.5$\% of A stars possess mid-infrared excesses at 22$\micron$, and the median lower limit to the fractional dust luminosity is $L_{\rm dust}/L_\ast\gtrsim 1.2\times10^{-6}$ for the A stars. At 12$\micron$, the occurrence rate of excesses is $0.08$\%$ \pm 0.04$\% for solar type stars and $1.0$\%$ \pm 0.5$\% for A stars. \item As a result of our study, the number of debris disks with known 10--30$\micron$ excesses within 75~pc (379) has now surpassed the number of disks with known $>$30$\micron$ excesses (289, with 171 in common), even if the latter are known to have a higher occurrence rate in unbiased samples. \end{enumerate} In addition to the scientific results, notable numerical and tabular references from the present study include: \begin{enumerate} \item the determination of photospherice \WS\ colors from $-0.15 < B_T-V_T < 1.4$~mag main sequecne stars (Table~\ref{tab:wise_bv_trends}) \item polynomial relations for correcting saturated \WS\ $4.5<W1<8.4$~mag and $2.8<W2<7.0$~mag photometry (Figure~\ref{fig:wise_flux_biases}) \item corrected $W1$ and $W2$ photometry for saturated \hip\ stars with $4.5<W1<8.4$~mag and $2.8<W2<7.0$~mag (Table~\ref{tab:stellarparameters}) \end{enumerate} \WS\ has rekindled the search for new disk bearing stars due to its enhanced resolving power compared to previous all-sky surveys like {\it IRAS}, combined with its wider coverage relative to pointed surveys using \spitzer. Although \WS\ cannot detect disks as faint as \spitzer, for that very reason the brighter, \WS-selected systems are excellent targets for resolved imaging observations, e.g., with the Gemini Planet Imager, ALMA, the LBTI nuller, or the JWST. Such observations would further constrain the structure of the disks and the properties of the dust grains that reside in them, expanding our knowledge of the range of planetary system architectures in the galaxy. | 14 | 3 | 1403.3435 |
1403 | 1403.4353_arXiv.txt | We report a detection of a faint near-Earth asteroid (NEA), which was done using our synthetic tracking technique and the CHIMERA instrument on the Palomar 200-inch telescope. This asteroid, with apparent magnitude of 23, was moving at 5.97 degrees per day and was detected at a signal-to-noise ratio (SNR) of 15 using 30 sec of data taken at a 16.7 Hz frame rate. The detection was confirmed by a second observation one hour later at the same SNR. The asteroid moved 7 arcseconds in sky over the 30 sec of integration time because of its high proper motion. The synthetic tracking using 16.7 Hz frames avoided the trailing loss suffered by conventional techniques relying on 30-sec exposure, which would degrade the surface brightness of image on CCD to an approximate magnitude of 25. This detection was a result of our 12-hour blind search conducted on the Palomar 200-inch telescope over two nights on September 11 and 12, 2013 scanning twice over six 5.0$^\circ\times$0.043$^\circ$ fields. The fact that we detected only one NEA, is consistent with Harris's estimation of the asteroid population distribution, which was used to predict the detection of 1--2 asteroids of absolute magnitude H=28--31 per night. The design of experiment, data analysis method, and algorithms for estimating astrometry are presented. We also demonstrate a milli-arcsecond astrometry using observations of two bright asteroids with the same system on Apr 3, 2013. Strategies of scheduling observations to detect small and fast-moving NEAs with the synthetic tracking technique are discussed. | \label{sec:intro} Recently, we introduced a technique of synthetic tracking that enabled detection of small and fast moving near-Earth asteroids (NEAs) using short exposure frames \citep{Shao2014}. Detecting and characterizing of small asteroids is important for several reasons. While being the subject of an interesting and rapidly-evolving area of planetary science, asteroids present a threat to the infrastructure and life on our planet. For example, a 17~m asteroid that hit Russia on Feb. 15, 2013, caused a severe damage to buildings and inflicted injuries to hundreds of people \citep{Brumfiel2013}. In addition, some NEAs may become targets for focussed space exploration efforts in the near future. In particular, characterizing small NEAs is needed to provide a potential target list for the upcoming NASA's asteroid redirection mission \citep{Lightfoot2013}. As discussed in \citep{Shao2014}, the synthetic tracking method processes the data from a number of short exposure frames by shifting each frame according to a tracking velocity vector so that the superposition of these shifted frames simulates a long-exposure integration with the telescope tracking at that velocity. This technique improves the detection's signal-to-noise ratio (SNR) by avoiding the trailing loss, which typically affects the detection of fast-moving NEAs at distances $\lesssim$ 0.1 AU from the Earth. This advantage is especially valuable for detecting small NEAs because these objects are observable only at short distances. The improved SNR from using the synthetic tracking technique yields a more precise astrometry of NEAs. In addition, synthetic tracking gains accruacy in astrometry from the reduction of the effects due to atmospheric disturbances and imprecise telescope pointing (to be addressed in Sec.~\ref{sec:astrometry}). Synthetic tracking technique is made possible by the availability of new generation cameras that provide both fast frame rate and low read noise. For example, scientific CMOS cameras can read at 100~Hz frame rate and only introduce read noise at 1e$^-$ level.\footnote{ \label{Andor2013} See a description of technical capabilities of the Andor's Neo and Zyla sCMOS Cameras at: {\tt http://www.andor.com/pdfs/literature/Andor\_sCMOS\_Brochure.pdf}} Our observation on Palomar 200-inch used the Andor's EMCCD$^{\ref{Andor2013}}$ operating at high EM gain of 200, making the read noise benign compared with the sky background noise even for the 16.7~Hz frame rate. This enabled us to detect a faint object at the apparent magnitude of 23 using about 30 sec of data ($\sim$ 500 frames). Without synthetic tracking, the detected asteroid with speed 5.97 degrees per day ($\d_p_d$) moves about 7 arcsecond ($''$) in the field. The corresponding surface brightness of the asteroid yields approximately ${\rm SNR}\approx 4$ (apparent magnitude 24.5), below the detection threshold of ${\rm SNR}=7$, set as our data processing criterium. An additional benefit of synthetic tracking is that it allows one to estimate velocity using only 30 sec of data, making confirmation task much easier. A 12-hour blind search was conducted on September 11 and 12, 2013, with six hours per night. Our survey continuously scanned over the sky at a rate of 5~$\as_p_s$, so that each star stays in the field for $\approx 30$~sec. It took about an hour to scan over each field of size 5.0$^\circ\times$0.043$^\circ$ (right ascention (RA) $\times$ declination (DEC)). We repeated the scan in the next hour to have two consecutive one-hour data covering the same field. Thus, during each night we covered three different 5.0$^\circ\times$0.043$^\circ$ fields with a total of six fields. the faint asteroid was detected in the second field on September 11, 2013, and confirmed by the repeated scan. Because this asteroid was observed twice, both times with SNR of $\sim$15, the false positive rate is practically zero. The asteroid moved 3770 pixels over 4626 seconds giving 5.97~$\d_p_d$ for the plate scale of 0.305~$''\!/$pixel. Using estimation of the asteroid population distribution from \citep{Harris2011}, we expected to detect 1--2 asteroids of H magnitude in range 28--31 per night. Our detection of one asteroid in the two-night survey is consistent with this expectation. The main step in the data processing with synthetic tracking technique is to search for signals (or bright spots) in the synthetically integrated images for a grid of tracking velocities. As mentioned in \citep{Shao2014}, this process is computationally intensive. To overcome this difficulty, we use graphics processing units (GPUs) to perform synthetic tracking at different velocities in parallel, thus, enable nearly real-time data processing. To improve the SNR, we adopt a matched filter scheme \citep{Turin1960} to low-pass filter the data using the point spread function (PSF) as the impulse response. For convenience, we use a Moffat's PSF function template \citep{Moffat1969} as our PSF model to quantify the shape of PSF and to reduce the computation in centroid fitting. We use a bright star in the field to determined PSF model parameters. A co-moving PSF fitting to the data frames is performed to generate precise astrometric solutions and velocities for both asteroids and the stellar objects present in a frame. The signal level resulting from the fitting is used to compute SNR and a false positive probability. This approach yields the same SNR as the matched filter scheme \citep{Gural2005, Shucker2008} in the case of using a template filter velocity matching that of the asteroid. The new feature in our approach is that we search in the tracking velocity space for the faint object using parallel computing and then optimize the tracking veloicty using the co-moving PSF fitting. In addition to the improvement of the detection SNR, synthetic tracking also yeilds more accurate astrometry for fast moving asteroids than the traditional long exposure approach. We achieved milli-arcsecond (mas) level astrometry for two known bright asteroids, observed on April 3, 2013, relative to nearby stars, after integrating over a minute using synthetic tracking. However, if using long exposures, as simulated by co-adding the short exposure frames, the astrometric precision does not improve after integrating over 30 sec. This is because the effect due to atmosphere and imprecise telescope pointing is no longer common between the asteroids and the background stars. Working on short exposure images, synthetic tracking technique makes the effect due to both the Earth's atmosphere and telescope pointing errors common between the asteroids and the background stars and thus achieves the similar precision of the relative astrometry for the asteroids to that of the relative stellar astrometry\citep{Boss2009}. This paper is organized as follows: In the Section~\ref{sec:syn_track}, we present the details of the synthetic tracking technique. The relevant algorithms are described in Section~\ref{sec:algorithm}. We present the data processing approach in Section~\ref{sec:data_processing}. In Section~\ref{sec:results}, we describe the results of the detection of the faint asteroid and also demonstrate precise astrometry results using data from observing two known asteroids. Observation strategies using synthetic tracking for detecting small NEAs are discussed. We conclude in Section~\ref{sec:conclude}. | \label{sec:conclude} We have presented the detection of a faint object as an application of the synthetic tracking technique to observational data taken on the Palomar 200-inch telescope including the data processing method and algorithms to demonstrate the efficacy. Synthetic tracking significantly improved detection SNR over traditional long exposure approach for fast moving NEAs, and thus enabled the detection of this faint object at apparent magnitude of 23. It also yields more precise astrometry by improving SNR and making the effects due to atmosphere and imprecise telescope pointing common between the asteroid and background stars, which cancels to the first order in relatively astrometry. Using the observational data of two known asteroids, we demonstrated milli aresecond level precision in astrometry. From a 12-hour blind search we found only one asteroid, which is consistent with Harris's population distribution of asteroids within the order of magnitude. Observation strategy using synthetic tracking to detect and discover faint objects within the short observation time window in general requires a well scheduled observation including a close to real time detection using 30 sec of observation data, confirmation within a few hours, and orbit determination within a week before the object becomes too faint to observe. | 14 | 3 | 1403.4353 |
1403 | 1403.3573_arXiv.txt | Since two years, the FACT telescope is operating on the Canary Island of La Palma. Apart from its purpose to serve as a monitoring facility for the brightest TeV blazars, it was built as a major step to establish solid state photon counters as detectors in Cherenkov astronomy. The camera of the First G-APD Cherenkov Telesope comprises 1440 Geiger-mode avalanche photo diodes (G-APD), equipped with solid light guides to increase the effective light collection area of each sensor. Since no sense-line is available, a special challenge is to keep the applied voltage stable although the current drawn by the G-APD depends on the flux of night-sky background photons significantly varying with ambient light conditions. Methods have been developed to keep the temperature and voltage dependent response of the G-APDs stable during operation. As a cross-check, dark count spectra with high statistics have been taken under different environmental conditions. In this presentation, the project, the developed methods and the experience from two years of operation of the first G-APD based camera in Cherenkov astronomy under changing environmental conditions will be presented. | \begin{figure}[!htb] \centering \includegraphics[width=0.49\textwidth]{DSCN5320.png} \caption{Picture of the FACT telescope during observations at night. In the background, the white container hosting the computing system, the bias power sully and the drive system is visible. Courtesy of D.~Dorner.} \label{fig:pic} \end{figure} For the first time in Cherenkov astronomy, photo multiplier tubes (PMT) for photo detection have been replaced by more modern silicon based photo sensors. The First G-APD Cherenkov Telescope (FACT) has implemented a focal plane using 1440 Hamamatsu MPPC S10362-33-050C sensors for photo detection. A picture of the telescope can be seen in Fig.~\ref{fig:pic}. A detailed description of its hardware and software is given in the FACT design report \cite{bib:design}. Apart from the use of recent technology for photo detection, the FACT project aims towards long-term monitoring of the brightest known blazars at TeV energies. Due to the natural gap in observations, which arises from bright moon light when operation with photo multiplier tubes is hardly possible with standard settings, and the busy schedule of existing high performance instruments, the existing sampling density is comparably low in the order of a few minutes every few days. Therefore, the aims is a high duty cycle instrument with the best possible data taking efficiency and observations during strong moon light, so that several hours per night and source can be achieved and the gap around full moon can be filled. Data taking efficiency and the results of almost two years of successful monitoring of Markarian 421 and Markarian 501 are presented at the end of the paper. In the following, the use of Geiger-mode avalanche photo sensors in the camera are discussed in more details and results from stability measurements are presented. \subsection{Application of G-APDs} In Cherenkov astronomy, photo sensors are used to detect Cherenkov light emitted by particle cascades in the atmosphere. The image of these particle showers is then used to reconstruct the properties of the primary particles such as particle type, energy and origin. To detect these nano-second short light-flashes consisting of a few tens to hundreds of photons, sensitive and fast photo sensors are needed. Since its beginning, Cherenkov astronomy has used photo multiplier tubes for photo detection, but for several years silicon based photo sensors are on the market as well. During the past years, Geiger-mode avalanche photo diodes (G-APD) became powerful and inexpensive enough to replace photo multiplier tubes in Cherenkov astronomy. Each sensor applied in the FACT camera consists of 3600 single Geiger-mode avalanche photo diodes (cells), 50\,\(\mu\)m\,x\,50\,\(\mu\)m each. Since the cells are operated above their breakdown voltage, the so-called Geiger-mode, each detected photon triggers a complete discharge of a single cell. The resulting pulse is comparable in amplitude and shape for all cells. Due to the high precision of the production, the fluctuations from cell to cell are small. Devices currently on the market achieve photo detection efficiencies comparable to the best available PMTs, and improved photo detection efficiency is expected for the next generation. As compared to PMTs, silicon based photo sensors are easier to handle due to their lower bias voltage in the order of 100\,V or below, they are mechanically more robust, and they are insensitive to magnetic fields. Apart from their high future potential, their main advantage in Cherenkov astronomy is their robustness against bright light which allows operation under bright moon light conditions. As a drawback, G-APDs encounter a high dark count rate, so-called optical crosstalk and a comparably high afterpulse probability. The dark count rate of the applied sensors is in the order of 1\,MHz to 9\,MHz for typical operation temperatures between 0\,\textdegree{}C to 30\,\textdegree{}C. At the same time, the rate of detected photons from the diffuse night sky background is in the order of 30\,MHz to 50\,MHz even during the darkest nights. Consequently, the dark count rate of the sensor is negligible. \subsection{Dark count rate, optical crosstalk and afterpulses} \begin{figure*}[!htb] \centering \includegraphics[width=0.48\textwidth]{crosstalk}\hfill \includegraphics[width=0.48\textwidth]{result} \caption{Left: The figure shows the relative error introduced on average by optical crosstalk (red) and its statistical error (black) on the signal as a function of the number of initial breakdowns considering a crosstalk probability of 11\%. In both cases, the error on a single event can be significantly larger, but on average the statistical error dominates considerably. Right: The figure shows the additional contribution from afterpulses (blue) as a fraction of the unbiased pulse as a function of the applied pulse integration window starting pulse integration at the half height leading edge compared to the contribution from optical crosstalk (red). While the relative contribution by afterpulses does not depend on the initial number of breakdowns, the relative contribution from optical crosstalk decreases with an increasing number of breakdowns. While the contribution from afterpulses in single events can reach 100\% the number of biased events is so small, that the average contribution for integration windows of less than 40\,ns is below 4\%.} \label{fig:result} \end{figure*} The so-called optical crosstalk is the probability that the discharge of a G-APD cell will indirectly trigger at least one other cell due to photons potentially emitted during the breakdown. While this effect is significant for signals in the order of one, for larger signals it just increases the amplitude and its fluctuation statistically. Figure~\ref{fig:result} shows the relative error introduced from crosstalk compared to the statistical error. Since the signals of Cherenkov showers are comprised of a number of photons large compared to one, and have an intrinsic fluctuation of the order of a few percent, optical crosstalk needs to be considered quantitatively in the shower reconstruction, but has no significant effect on the quality. Afterpulses in G-APDs originate from trapped charges released with a short delay after the initial pulse. Their probability is exponentially decreasing with a time scale in the order of the decay time of the pulse, see \cite{bib:afterpulses}. Figure~\ref{fig:result} shows the influence of afterpulses as a function of the integration window. Although their total probability in the order of 10\% to 20\% is high compared with PMTs, they usually occur within the timescale of the initial pulse. The exponential decrease of the probability ensures that afterpulses induced by coincident signals in several cells or sensors are not released synchronously. Therefore, afterpulses in G-APDs can not fake trigger signals in contrast to afterpulses in PMTs. In addition, due to the partial discharge of the cell after the initial pulse, the amplitude of early afterpulses is highly decreased. Generally, their absolute amplitude, i.e. falling edge plus afterpulse, does not exceed the amplitude of the initial pulse. In the pulse extraction algorithm, their influence can be suppressed strongly, if the integration range is limited to a short interval around the peak of the initial pulse. \subsection{Breakdown voltage} The main challenge in the application of G-APDs in an experiment exposed to changing environmental conditions is the dependence of their breakdown voltage from the temperature and the voltage drop induced by high currents. Cherenkov telescopes are exposed to natural temperature changes during the night. Since the breakdown voltage of the G-APDs in use is temperature dependent (\(\approx\)55\,mV/K), the gain will follow the temperature, if the voltage is not adjusted accordingly. In addition, each sensor has a network of serial resistors. When the moon rises during the night, more breakdowns take place due to the increased photon flux and induce a higher current in the resistors. The voltage drop induced at the serial resistors is enough to significantly decrease the gain of the sensor. Therefore, this voltage drop must be determined and the applied voltage has to be adjusted accordingly. In the following, the calibration and adjustment procedure applied in the FACT camera is presented and discussed. | The First G-APD Cherenkov telescope has proven the possibility to apply Geiger-mode avalanche photo diodes in Cherenkov astronomy for photo detection. Their gain has been shown to be independent of the temperature within 1\% in a temperature range between 0\,\textdegree{}C to 30\,\textdegree{}C and independent from the night-sky background light level within at least 6.5\%. The gain distribution of the 1440 pixels is not wider than 3.4\% with a closed lid and 10\% with open lid, which includes the inhomogeneity of the light-pulser's light distribution used for the measurement. These stability has been achieved independent of any external calibration device. The external light-pulser has only be used to measure the G-APD properties. Currently, more calibration measurements are taken to improve these results further. With a high level of automation, the fully remote controlled telescope, achieves a data taking efficiency of typical 85\% and more than 95\% if calibration measurement are included. Configuration and re-positioning times are less than 5\% in total. Recently, it was proven that observations even close to the full moon are possible, c.f.~\cite{bib:moon}. Monitoring of the two brightest known blazars is carried out successfully since two years and the excess rates of the past 1.5 years have been present including a couple of minor and major flares. The results from an online analysis are available online immediately after data taking. \vfill | 14 | 3 | 1403.3573 |
1403 | 1403.3090_arXiv.txt | {An increasing number of high-resolution stellar spectra is available today thanks to many past and ongoing spectroscopic surveys. Consequently, numerous methods have been developed to perform an automatic spectral analysis on a massive amount of data. When reviewing published results, biases arise and they need to be addressed and minimized.} {We are providing a homogeneous library with a common set of calibration stars (known as the Gaia FGK benchmark stars) that will allow us to assess stellar analysis methods and calibrate spectroscopic surveys.} {High-resolution and signal-to-noise spectra were compiled from different instruments. We developed an automatic process to homogenize the observed data and assess the quality of the resulting library.} {We built a high-quality library that will facilitate the assessment of spectral analyses and the calibration of present and future spectroscopic surveys. The automation of the process minimizes the human subjectivity and ensures reproducibility. Additionally, it allows us to quickly adapt the library to specific needs that can arise from future spectroscopic analyses. } {} | Investigations into how the Milky Way is formed and its evolution are being revolutionized thanks to the many ongoing stellar spectroscopic surveys such as SDSS \citep{2000AJ....120.1579Y}, LAMOST \citep{2006ChJAA...6..265Z}, RAVE \citep{2006AJ....132.1645S}, Gaia \citep{2001A&A...369..339P}, Gaia-ESO \citep[GES,][]{2012Msngr.147...25G}, HERMES/GALAH \citep{2010gama.conf..319F} and APOGEE \citep{2008AN....329.1018A}. Tracing the chemical and dynamical signatures of large samples of stars helps us to distinguish the different Galactic components and thus understand when and how the different Galactic formation scenarios took place. The quantity of spectroscopic data available today requires the development of automatic spectral analysis. Numerous methods have been developed over the past years \citep[e.g.,,][to name a few]{1996A&AS..118..595V, 1998A&A...338..151K, 2006MNRAS.370..141R, 2008AJ....136.2022L, 2009A&A...501.1269K, 2010A&A...517A..57J, 2012A&A...544A.154P, 2013ApJ...766...78M, 2013A&A...558A..38M} to asses large datasets, where each of them have different approaches to calibrate and evaluate their results. However, each survey has its own setup (e.g., spectral range, resolution) and each spectral analysis code has its own particularities (i.e., continuum normalization, atomic line lists). The consequence is that the resulting parameters cannot be directly combined and used for galactic and stellar studies. Thus, spectroscopic calibration with a common reference set of stars is required. There are several stellar spectral libraries available in the community that are used for calibration in some sense \cite[see, e.g.,][for a compilation]{2005MSAIS...8..170M}, providing a large sample of good spectra of stars covering a large part of the Hertzsprung-Russel (HR) diagram and metallicities. Examples of them are ELODIE \citep{2001A&A...369.1048P}, Indo-US \citep{2004ApJS..152..251V}, MILES \citep{2006MNRAS.371..703S}, StarCAT \citep{2010ApJS..187..149A}, and UVES-POP \citep{2003Msngr.114...10B}. These libraries contain a large number of stars (usually above 1,000) and they differ from each other in terms of resolution and wavelength coverage. They are frequently used for stellar population synthesis models and galactic studies \citep[e.g.,][and references therein]{2012MNRAS.424..157V, 2009ApJ...690..427P, 2005MNRAS.364..503Z} and for calibration or validation of methods that determine stellar parameters from stellar spectra \citep[e.g][]{2008AJ....136.2070A, 2009A&A...501.1269K, 2011RAA....11..924W}. Nevertheless, the Sun is frequently the only calibration star in common between different methods/surveys and, depending on the survey, its observation is not always possible. Our motivation for defining the Gaia FGK benchmark stars is to provide a common set of calibration stars beyond the Sun, covering different regions of the HR diagram and spanning a wide range in metallicity. They will be used as pillars for the calibration of the parameters that will be derived for one billion stars by Gaia \citep{2001A&A...369..339P}. The defining property of these stars is that we know their radius and bolometric flux, which allows us to estimate their effective temperature and surface gravity {\it fundamentally}, namely, independent of the spectra. In Heiter et al (in prep, Paper~I), we provide the main properties of our sample of the Gaia FGK benchmark stars and describe the determination of temperature and gravity. In this article (Paper~II), we introduce the spectral library of the Gaia FGK benchmark stars. In \citet[][Paper~III]{2013arXiv1309.1099J}, we analyse our library with the aim to provide a homogeneous scale for the metallicity. The current sample of Gaia benchmark stars is composed of bright, well-known FGK dwarfs, subgiants, and giants with metallicities between solar and $-2.7$~dex. We selected stars for which angular diameter and bolometric flux measurements are available or possible. They have accurate parallax measurements, mostly from the HIPPARCOS mission. The sample contains several visual binary stars. In particular, both the A and B components are included for the $\alpha$ Cen and the 61 Cyg systems. The star $\eta$ Boo is a single-lined spectroscopic binary \citep{2005A&A...436..253T}. The fastest rotators in the sample are $\eta$ Boo and HD 49933, with $v\sin i \gtrsim 10$~km~s$^{-1}$. The metal-poor dwarf Gmb 1830 has the highest proper motion (4.0 and $-$5.8 arcsec/yr in right ascension and declination, respectively). Most of the other stars have proper motions less than 1 arcsec/yr. Our library provides a homogeneous set of high-resolution and high signal-to-noise ratio (S/N) spectra for the 34 benchmark stars. Moreover, the stellar parameters of these benchmark stars were determined consistently and homogeneously, making them perfect for being used as reference. This library of 34 benchmark stars is therefore a powerful tool to cross-calibrate methods and stellar surveys, which is crucial for having a better understanding of the structure and evolution of the Milky Way. The observed spectra of the benchmark stars were obtained from different telescopes with different instruments and specifications (i.e., resolution and sampling). We developed an automatic process to transform the spectra into one final homogeneous dataset. This allows us to easily generate new versions of the library adapted to the needs of specific surveys (i.e., downgrading the resolution or selecting a different spectral region). Additionally, since reproducibility is one of the main pillars of science, our code will be provided under an open source license to any third party wishes to reproduce the results \citep{2013arXiv1312.4545B}. This article is structured as follows. In Sect.~\ref{sub:observational_data}, we describe the original observed spectra and its sources, while in Sect.~\ref{sub:data_handling_and_processing} we introduce the computer process that was developed to create the library. Section~\ref{sub:validation} presents the different tests that were performed to validate the correctness of processing and the consistency of the library. In Sect.~\ref{sub:results}, we describe the resulting library's elements that we provide, and finally, we conclude the paper in Sect.~\ref{sub:conclusions} | \label{sub:conclusions} We created a homogeneous library of high resolution and high S/N spectra corresponding to 34 benchmark stars with four different variants (convolved/not convolved original non-normalized fluxes, convolved/not convolved normalized spectra). The library provides a powerful tool to assess spectral analysis methods and calibrate spectroscopic surveys. We validated the consistency of the library by carefully checking the normalization and convolution treatments. The radial velocity corrections was certified by comparing the results with the high precision measurements of HARPS pipeline. We verified the coherence of the treated spectra by comparing them with EW measurements completely independent from our process. These strict tests proved the high quality level of the spectral library. The whole creation and verification process was automatized, minimizing human subjectivity and ensuring reproducibility. It also allows us to create new versions of the library adapted to particular needs (i.e., different resolutions and spectral ranges) of specific spectroscopic surveys or spectral analyses. The Gaia FGK benchmark stars library provides an opportunity to homogenize spectroscopic results (from single observations to massive surveys), reducing their dispersion and making them more comparable. This higher level of homogeneity can lead to a better and more robust understanding of the Galaxy such as its formation, evolution, and current structure. | 14 | 3 | 1403.3090 |
1403 | 1403.1576_arXiv.txt | It is pointed out that oscillating current density, produced due to the coupling between an external magnetic field and the cosmic axion field, can excite the TM resonant modes inside an open-ended cavity (tube). By systematically solving the field equations of axion-electrodynamics we obtain explicit expressions for the oscillating fields induced inside a cylindrical tube. We calculate the enhancement factor when a resonance condition is met. While the power obtained for TM modes replicates the previous result, we emphasize that the knowledge of explicit field configurations inside a tube opens up new ways to design axion experiments including a recent proposal to detect the induced fields using a superconducting LC circuit. In addition, as an example, we estimate the induced fields in a cylindrical tube in the presence of a static uniform magnetic field applied only to a part of its volume. | 14 | 3 | 1403.1576 |
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1403 | 1403.2605_arXiv.txt | {} {We study the impact of binary interaction processes on the evolution of low- and intermediate-mass stars using long-term monitoring of their radial velocity. Here we report on our results on the central stars of two planetary nebulae (PNe): the well-studied spectrophotometric standard BD+33$^{\circ}$2642 (central star of PNG 052.7+50.7) and HD\,112313 (central star of PN LoTr5), the optical light of which is dominated by a rapidly rotating G star.} {The high-resolution spectra were cross-correlated with carefully selected masks of spectral lines. The individual masks were optimised for the spectral signatures of the dominant contributor of the optical light.} {We report on the first detection of orbital motion in these two objects. For \bd\, we sampled 1.5 cycles of the 1105 $\pm$ 24 day orbital period. For HD\,112313 a full period is not yet covered, despite our 1807 days of monitoring. The radial-velocity amplitude shows that it is unlikely that the orbital plane is co-planar with the one defined by the nebular waist of the bipolar nebula. To our knowledge these are the first detections of orbits in PNe that are in a range from several weeks to a few years.} {The orbital properties and chemical composition of \bd\, are similar to what is found in post-AGB binaries with circumbinary discs. The latter are probably progenitors of these PNe. For LoTr5 the Ba-rich central star and the long orbital period are similar to the Ba star giants, which hence serve as natural progeny. In contrast to the central star in LoTr5, normal Ba stars are slow rotators. The orbits of these systems have a low probability of occurrence according to recent population synthesis calculations. } | The role of binary interaction processes in the shape and shaping of PNe and their progenitors, the proto-planetary nebulae (PPNe), is far from understood \citep[e.g.][and references therein]{balick02,vanwinckel03,demarco09}. There is growing consensus that binary interaction processes play a fundamental role in many objects, but the theoretical understanding is subject to many unsupported assumptions such as the efficiency of envelope ejection, the physical description of the common-envelope phase (CE phase, \citet{izzard12,ivanova13}), the accretion efficiency onto the companion \citep{ricker08}, and the jet formation mechanisms in binaries. It also is an observational challenge to constrain the requirements for the most important shaping agents such as jets or magnetic fields and determine whether they are active or not. Direct observational evidence of the binary nature of the central star of PNe and obscured PPNe is notoriously difficult to obtain \citep{demarco09} and all the techniques have their specific biases. The most successful method so far is the detection of orbital modulation in the light curve \citep[e.g.][]{bond00b,demarco08,miszalski09a,miszalski09b,miszalski11b}, which led to the discovery of about 40 systems. The modulation can be caused by ellipsoidal deformation, eclipses, or irradiation of the cool component by the hot component \citep[e.g.][]{hillwig10}. These photometric methods are only sensitive to systems that have spiraled-in during a CE phase and have now a period shorter than about 10 days. Several of these post-CE systems also show jets powered by accretion \citep[e.g.][]{boffin12, miszalski13a}. It is as yet unclear under which conditions wider systems will create accretion-powered jets; the main reason is the lack of observational data on the orbital characteristics of the central stars. Here we report on the first discovery of two wide spectroscopic binaries among central stars of PNe. We introduce the objects (Sect.~\ref{sect:objects}) and our radial-velocity programme (Sect~\ref{sect:radvel}) in the next sections. We focus on the results in Sect.~\ref{sect:bd} and Sect.~\ref{sect:LoTr5} and conclude in Sect.~\ref{sect:results}. | \begin{itemize} \item \bd\, is a spectrophotometric standard star and one of the few PNe in the Galactic halo. It is a binary with abundance anomalies, very similar to what is found in many Galactic post-AGB binaries with circumbinary discs. The orbit of \bd\, is similar to the orbits found in the post-AGB binaries. This object is therefore a natural progeny of the class of depleted F-K type post-AGB binaries with circumstellar discs \citep{vanwinckel03} and shows that the latter can also become PNe. \item Although the full orbit of HD\,112313 is not yet covered, the lower limit of the mass function led to a conflict between the nebular and orbital properties. The obtained mass function implies that either the orbital plane is not co-planar with the waist of the bipolar nebula, or that the WD is member of a close binary with a total mass of 3.5 M$_{\odot}$. The cool Ba-rich component of the system must be spun-up by wind accretion, and we disproved the statements in the literature that the rapidly rotating G star may be a close binary. \end{itemize} Our project shows that detecting radial-velocity orbits in central stars of PNe is possible but intensive long-term monitoring is clearly needed. This requires a dedicated long-term observational project. Many prime candidates for PNe with binary central stars can be found, but they are very faint \citep[e.g.][]{demarco13} and will require regular telescope time on large infrastructure. | 14 | 3 | 1403.2605 |
1403 | 1403.7202_arXiv.txt | We present follow-up optical imaging and spectroscopy of one of the light echoes of $\eta$~Carinae's 19th-century Great Eruption discovered by \cite{Rest12a}. By obtaining images and spectra at the same light echo position between 2011 and 2014, we follow the evolution of the Great Eruption on a three-year timescale. We find remarkable changes in the photometric and spectroscopic evolution of the echo light. The $i$-band light curve shows a decline of $\sim 0.9$~mag in $\sim 1$~year after the peak observed in early 2011 and a flattening at later times. The spectra show a pure-absorption early G-type stellar spectrum at peak, but a few months after peak the lines of the \ion{Ca}{2} triplet develop strong P-Cygni profiles and we see the appearance of [\ion{Ca}{2}]~7291,7324 doublet in emission. These emission features and their evolution in time resemble those observed in the spectra of some Type~IIn supernovae and supernova impostors. Most surprisingly, starting \mbox{$\sim 300$~days} after peak brightness, the spectra show strong molecular transitions of CN at $\gtrsim 6800$~\AA. The appearance of these CN features can be explained if the ejecta are strongly Nitrogen enhanced, as is observed in modern spectroscopic studies of the bipolar Homunculus nebula. Given the spectroscopic evolution of the light echo, velocities of the main features, and detection of strong CN, we are likely seeing ejecta that contributes directly to the Homunculus nebula. | $\eta$~Carinae ($\eta$~Car) is one of the most well observed and, at the same time, mysterious and poorly understood objects in the night sky \citep[e.g.,][and references therein]{DHbook}. At a distance of $d \simeq 2.35$~kpc \citep{Smith06} in the young cluster Trumpler~16 in the Carina nebula, $\eta$~Car is a very luminous and massive evolved star ($L\sim 10^{6.7}$~L$_\odot$, $M\sim 100$~M$_\odot$; e.g., \citealt{Hillier2001}), a high mass-loss Luminous Blue Variable (LBV), in an eccentric binary system \citep{Damineli96}. The system is surrounded by a dusty, massive ($\sim 20$~M$_\odot$, \citealt{Smith03}) bipolar nebula, the Homunculus nebula, that was ejected $\sim 170$~years ago in the Great Eruption \citep{Currie96,Smith98,Morse2001}. The Great Eruption (GE) of $\eta$~Car in the mid-1800s ($\sim 1840-1860$) was a spectacular astronomical event, a bright and energetic transient visible to the naked eye that was observed and registered by many \citep[e.g.,][]{Herschel1847}. Although we do not understand the physical mechanism that caused $\eta$~Car's GE, it is used as a standard reference for understanding episodic mass-loss in very massive stars \citep[e.g.,][]{Smith14}, LBVs \citep[e.g.,][]{Humphreys94}, supernova impostors \citep[e.g.,][]{Humphreys99,vandyk02,Maund06,Smith11,Kochanek2012}, and the most luminous supernovae \citep[e.g.,][]{Smith07}. Yet, until recently we only knew it through visual estimates of its brightness and color, due to the technological limitations at the time of this event \citep{Frew04,SF11}. The recent discovery of $\eta$~Car's light echoes (LEs) by \citet[][hereafter R12a]{Rest12a} now gives us the opportunity to re-observe the multiple brightening events of $\eta$~Car in the mid-1800s with modern instrumentation and from multiple directions \citep{Rest12b}, similar to what has been done for ancient supernovae in the LMC \citep{Rest05b,Rest08a}, SN~1987A \citep{Sinnott13}, Cas~A \citep{Rest08b,Krause08a, Rest11_casaspec,Rest11_leprofile}, and Tycho \citep{Rest08b,Krause08b}. In R12a we presented initial imaging and spectroscopy of one of the brightest LEs of $\eta$~Car which showed its unambiguous association with one of the brightenings during the GE instead of the smaller eruption of circa 1890 \citep{Humphreys99}. The spectra of this LE obtained in early 2011 showed pure absorption lines consistent with supergiants with spectral type G2$-$G5 ($\rm T_{eff} \sim 5000$~K) and line velocities of $\sim -200$~km~s$^{-1}$. In this letter, we present three years of imaging and spectroscopic follow-up observations of the light echo of $\eta$~Car's GE discovered by R12a. In Section~\S\ref{sec2} we discuss the observations and data reduction. In Section~\S\ref{sec3} we present the results and analysis of the light curve and spectra. We discuss our results in Section~\S\ref{sec4}. | \label{sec4} We have presented three years of photometric and spectroscopic follow-up observations, from early 2011 to 2014, of the light echo of $\eta$~Car's 19th century Great Eruption discovered by R12a. These observations give us a unique opportunity to study the evolution and physical properties of the GE in detail. The detection of P-Cygni profiles and the observed velocities of the spectral features unambiguously link $\eta$~Car's GE to luminous extragalactic transients associated with episodic mass-loss events in massive stars, as has been suggested in previous studies based on its historical light curve and the properties of the Homunculus nebula \citep[e.g.,][]{Humphreys99,vandyk02,Smith11,Smith13}. Indeed, the appearance and time evolution of strong \ion{Ca}{2} triplet, from pure absorption features at peak to P-Cygni profiles to emission-dominated profiles at late times, and [\ion{Ca}{2}] forbidden doublet in the spectra of $\eta$~Car's light echo resemble the spectral properties and evolution of some Type~IIn supernovae (SN~1994W, \citealt{Dessart09}; SN~2011ht, \citealt{Mauerhan13}) and supernova impostors \citep[e.g., UGC~2773-OT,][]{Smith10}. The time evolution of the line profiles of the LE is broadly consistent with an optical depth effect in the ejecta, as seen in Type~II SN. We can measure the electron density in the \ion{Ca}{2} line-forming region using the ratio of intensities of the doublet to the triplet \citep{FerlandPersson89}. From the light echo spectra, we measure an intensity ratio of $\sim 1$, roughly constant at different epochs. This gives an electron density in the range $\sim 10^8-10^{10}$~cm$^{-3}$ for $\rm T_e \simeq 3000-20000$~K. This high density is consistent with modeling of the spectra of the Type~IIn SN~1994W \citep{Dessart09}, which also shows a similar spectral evolution in the \ion{Ca}{2} lines. However, we do not see in these LE spectra the broad H$\alpha$ wings produced by electron scattering in Type~IIn SN. The comparison to the SN impostor UGC~2773-OT is particularly relevant here because the FWHM of the H$\alpha$ line ($\sim 400$~km~s$^{-1}$) and its relative strength to the \ion{Ca}{2} emission \citep{Smith10} are consistent with the non-detection of broad H$\alpha$ features in $\eta$~Car's light echo spectra. Despite the similarities between some of the properties of $\eta$~Car's LE spectra and SN impostors, there are important and interesting differences. As pointed out in R12a, the LE spectrum at peak is consistent with cool supergiants with late spectral type G2$-$G5 ($\rm T_{eff} \sim 5000$~K). This is cooler than all the published spectral analysis studies of SN impostors and LBV outbursts to date\footnote{For example, the coolest well-measured temperature of an LBV in outburst is that of R71 in the LMC, which had $\rm T_{eff} \sim 6650$~K in 2012 \citep{mehner13}.}. The new, late-time spectra of the LE (starting in the spectrum obtained $\sim 300$~days after peak) confirm this result. Spectral fitting of the blue part of the spectra using UlySS and cross-correlation with observed spectra of supergiants give $\rm T_{eff}\sim 4000-4500$~K (Rest et al. 2014, in prep.). Furthermore, we detect strong CN molecular bands in the red part of the optical spectra, which are seen in the spectra of cool stars (e.g., carbon stars and RSGs) but are not expected to be formed in warmer stellar photospheres with $\rm T_{eff} \gtrsim 5500$~K \citep[e.g.,][]{Lancon2007}. This indicates that the spectrum of the LE likely becomes cooler after the light curve peak and stays cool at late times. The cool temperature of the LE at peak and at late times after peak are inconsistent with the most simple predictions from the opaque wind model for the GE \citep{Davidson87}, and also very different from S-Doradus variations and some LBV giant eruptions \citep[e.g.,][]{Smith11}. The model of the GE as an explosion and interaction with a strong pre-GE wind \citep{Smith13} appears to be more consistent with our observations. However, we note that further analysis and spectra of multiple light echoes as a function of time will be needed in order to directly test different GE models \citep[e.g.,][]{Davidson87,Soker07,Smith13}. The CN features detected in the LE spectra at late times are significantly stronger than expected from stellar atmosphere models for RSGs (see Figure~\ref{fig5}). We could not find in the literature any observed spectra with stronger features, relative to the continuum, even among the $\sim 1000$ carbon stars identified in SDSS (G.~Knapp, private communication) and also among different types of massive star explosions and outbursts (SNe and SN impostors). This points to an abundance effect, in the sense that the material ejected in the GE is likely significantly rich in Nitrogen. This is consistent with the measured over-abundance of N in the Homunculus nebula, derived from ionized \citep[e.g.,][]{Davidson82,Dufour97} and from molecular species \citep[e.g.,][]{Smith06b,Loinard12}. The spectroscopic evolution of the LE, the observed velocities of the \ion{Ca}{2} features (which are consistent with the velocities of the Homunculus at similar latitudes presented in \citealt{Smith06}), and the strong CN features pointing to N enhancement, suggest that we are seeing ejecta that contributes directly to the Homunculus nebula. | 14 | 3 | 1403.7202 |
1403 | 1403.0939_arXiv.txt | Understanding how black holes accrete and supply feedback to their environment is one of the outstanding challenges of modern astrophysics. \source\ is a candidate black hole low-mass X-ray binary that was discovered in 2012 when it entered an accretion outburst. To investigate the binary configuration and the accretion morphology we monitored the evolution of the outburst for $\simeq$3 months at X-ray, UV, optical ($B,V,R,I$), and near-infrared ($J,H,K$) wavelengths using \swift\ and \smarts. The source evolved from a hard to a soft X-ray spectral state with a relatively cold accretion disk that peaked at $\simeq$0.5~keV. A \chan/HETG spectrum obtained during this soft state did not reveal signatures of an ionized disk wind. Both the low disk temperature and the absence of a detectable wind could indicate that the system is viewed at relatively low inclination. The multi-wavelength light curves revealed two notable features that appear to be related to X-ray state changes. Firstly, a prominent flux decrease was observed in all wavebands $\simeq1-2$ weeks before the source entered the soft state. This dip occurred in (0.6--10 keV) X-rays $\simeq6$~days later than at longer wavelengths, which could possibly reflect the viscous time scale of the disk. Secondly, about two weeks after the source transitioned back into the hard state, the UV emission significantly increased while the X-rays steadily decayed. We discuss how these observations may reflect changes in the accretion morphology, perhaps related to the quenching/launch of a jet or the collapse/recovery of a hot flow. | Black holes in low-mass X-ray binaries (LMXBs) accrete matter from a low-mass ($\lesssim$1$~\Msun$) companion star that overflows its Roche lobe. During accretion outbursts, matter is supplied to the black hole through an accretion disk and outflows in the form of disk winds and jets are generated. X-ray monitoring has revealed a rich timing and spectral behavior leading to the formulation of distinct X-ray spectral states and transitions between them \citep[e.g.,][]{homan2005_specstates,remillard2006}.\footnote[10]{In this work we discuss X-ray spectral data in the 0.6--10 keV energy range and therefore we adopt a simplified soft/hard state nomenclature that neglects intermediate states.} During {\it soft} states, the X-ray spectrum is dominated by thermal emission from the hot inner part of the accretion disk with a typical temperature of $kT_{\mathrm{in}}\simeq$1~keV \citep[e.g.,][]{dunn2011,reynolds2013}. Strong ionized X-ray disk winds are commonly detected during this state \citep[e.g.,][]{miller2006_winds,ponti2012_winds}. In {\it hard} states, the accretion disk is colder and the X-ray spectrum is dominated by a power law with a photon index of $\Gamma \simeq 1.5-2$. Disk winds seem to be suppressed \citep[e.g.,][]{miller2008,miller2012_winds,neilsen2009}, and instead compact radio jets are observed \citep[e.g.,][]{fender2005}. The hard X-rays are ascribed to a hot disk corona, a radiatively inefficient accretion flow, or a jet \citep[e.g.,][]{esin1997,markoff2001,brocksopp2004}. There are several mechanisms that are thought to contribute to the optical, ultraviolet (UV) and near-infrared (nIR) emission of black hole LMXBs. This radiation may originate in the cool outer parts of the accretion disk as the result of viscous dissipation or reprocessing of X-rays, but could also be produced in the jet, corona or inner hot flow \citep[e.g.,][]{paradijs95,esin1997,markoff2001,russell2006,rykoff2007,veledina2011}. The late-type companion star is typically dim and not contributing significantly to the outburst flux. Multi-wavelength observations of the outburst evolution can shed light on the physical mechanisms producing the different emission components, the accretion inflow-outflow coupling, and the morphological changes occurring during X-ray state transitions. In this work we report on such a study of the newly discovered candidate black hole LMXB \source. | \label{sec:discussion} \subsection{Inclination and Disk Temperature} A \chan/HETG spectrum obtained when \source\ was in the soft state appeared featureless, despite the fact that ionized disk winds are expected to be ubiquitous in this spectral state \citep[e.g.,][]{miller2008,miller2012_winds}. If such winds are mostly equatorial, a possible explanation for the lack of ionized absorption features is that \source\ is viewed at relatively low inclination \citep[][]{ponti2012_winds}. This would be consistent with the analysis of disk reflection features, which suggested an inclination of $i \lesssim 20^{\circ}$ \citep[][]{reis2013}. Geometrical effects could also account for the fact that the accretion disk of \source\ peaked at $kT_{\mathrm{in}}\simeq$0.5~keV, which is much cooler than typically seen for black hole LMXBs in the soft state \citep[$kT_{\mathrm{in}}\gtrsim$1~keV; e.g.,][]{dunn2011,reynolds2013}. As shown by \citet{munozdarias2013}, disks look cooler when viewed at lower inclination, which can be understood in terms of relativistic effects. \source\ could thus fit into this picture. Alternatively, the low temperature could be the result of a small disk size, e.g., due to a short orbital period (see Section~\ref{subsec:visc}), a truncated disk or a retrograde black hole spin \citep[][]{reis2013}. \subsection{Viscous Time Scale of the Accretion Disk?}\label{subsec:visc} Our multi-wavelength light curves revealed a prominent flux dip that first appeared in the optical band, only to be followed by X-rays $\simeq$6~days later. A comparable X-ray lag ($\simeq$8~days) is suggested by cross-correlating the light curves over the full time range covered by our monitoring campaign. A natural explanation for the observed delay may be a (small-scale) mass-transfer instability that originated at the outer edge of the accretion disk and then propagated inward. The time delay could then represent the viscous time scale of the disk, $t_{\mathrm{visc}}(r) = \frac{2}{3\alpha}(\frac{h}{r})^{-2}\Omega_{K}^{-1}$. Here $\alpha$ is the viscosity parameter, $h$ the scale height of the disk, and $\Omega_{K}$ the Keplerian frequency at radius $r$. Assuming $\alpha=0.1$, $h/r=0.01$ (i.e., a geometrically thin disk), and $M=8~\Msun$, a viscous time scale of $t_{\mathrm{visc}}=6$~days would yield a disk radius of $r\simeq3400~GM/c^2$ ($\simeq4\times10^9$~cm). This is relatively small, perhaps supporting the fairly short $\simeq$2--4 hr orbital period hinted by its rapid optical flaring (\cite{lloyd2012}, but see \cite{casares2012}). On the other hand, a comparable time delay of several days between the nIR and X-rays has also been inferred for black hole LMXBs that have (much) longer orbital periods \citep[e.g., LMC--X3, GX 339--4, and 4U 1957+11;][]{brocksopp2001,homan2005,russell2010}. The interpretation of the dip as a mass-transfer instability could naturally explain why in all bands the flux recovered to the pre-dip level. However, this picture may not account for the fact that the flux appeared to hit a minimum in the $H$ and $K$ bands later than in the optical/UV (albeit still before the X-rays; Figure~\ref{fig:dip}). This may suggest that multiple (competing) emission mechanisms are operating, or that the flux dip is the result of a different physical process altogether (Section~\ref{subsec:flow}). This is perhaps hinted by the fact that there appears to be a causal connection between the dip and the transition to the soft state $\simeq1-2$~weeks later. This connection would not be obvious to explain in terms of a small-scale mass-transfer instability. \subsection{Changes in the Accretion Flow Morphology}\label{subsec:flow} Interestingly, a very similar flux dip as we observed for \source\ was reported for \gx\ during its 2010 outburst \citep[][]{yan2012}. In that source the UV flux decreased by $\simeq$60\% in $\simeq$2 weeks time, which was associated with a drop in the optical/nIR (by $\simeq$85\%) as well as the radio flux. Within $\simeq$10 days after this drop the source transitioned from a hard to soft state, and the UV flux dip was therefore associated with jet quenching \citep[][]{yan2012}. In \source, support for change in accretion morphology (perhaps connected to a jet) is provided the observation that the nIR color suddenly became bluer right around the time of the flux dip (Figure~\ref{fig:lc}). It remained as such until the end of the soft state, which started $\simeq10$~days after the color change. This behavior was not observed in the optical bands, suggesting a relative suppression of the nIR emission during the soft state, i.e., when jets are known to be quenched. When the source transitioned back to the hard state, the $K$ and $H$ band flux started to rise and the nIR color become redder again (around day 138 in Figure~\ref{fig:lc}). Although our \smarts\ monitoring stopped at this time, the hinted nIR brightening may be similar to that seen in other black hole LMXBs entering the hard state \citep[e.g.,][]{jain2001,buxton2004,buxton2012,kalemci2005,kalemci2013,russell2010,russell2013,dincer2012,corbel2013}. The optical/nIR secondary maxima seen in these sources have been ascribed to the revival of a jet \citep[e.g.,][]{dincer2012,kalemci2013}. Continued \swift\ monitoring revealed that in \source\ the UV emission started to increase $\simeq$20 days after the transition back into the hard state (around day 158 in Figure~\ref{fig:lc}), whereas the X-rays continuously decayed. This caused the relation between the UV and 0.6--10 keV flux to move from a positive correlation (with a slope of $\beta\simeq0.15$) into a negative correlation ($\beta\simeq -0.45$). This turnover in the UV/X-ray relation occurred at $F_X\simeq10^{-9}~\flux$ (0.6--10 keV). The distance toward \source\ is unknown, but if the peak of the outburst (as observed by \maxi) did not exceed the Eddington limit ($L_{\mathrm{EDD}}$), the transition flux corresponds to $\lesssim0.03\times L_{\mathrm{EDD}}$. The UV brightening was accompanied by a dramatic increase in the X-ray hardness ratio, which again suggests a causal connection with changes in the accretion morphology. This rise in UV emission is reminiscent of the secondary nIR/optical maxima seen in other black hole LMXBs. In a systematic study of a number of different sources and outbursts, \citet{kalemci2013} showed that the nIR peaks $\simeq5-15$~days after the transition from the soft to the hard state. In \source, the UV appears to peak $\simeq10$ days after it starts rising, or $\simeq30$~days after the state transition (around day 168 in Figure~\ref{fig:lc}). This is thus considerably later than the reported nIR peaks \citep[][]{kalemci2013}. In a jet interpretation it might be expected that the UV increases after the nIR \citep[][]{corbel2013,russell2013}, but it is not immediately clear if this delay could be as large as $\simeq20$~days. An alternative interpretation for both the multi-wavelength flux dip and the UV peak may be provided in terms of the collapse and recovery of a hot inner flow \citep[][]{veledina2011,veledina2013}. Such a radiatively inefficient flow might replace the inner accretion disk at low mass-accretion rates \citep[e.g.,][]{esin1997}. Analysis of disk reflection features suggested that the disk in \source\ could indeed be subject to evaporation \citep{reis2013}. This model predicts that the outer regions of the hot flow (traced by the optical/nIR emission) collapses several days before an X-ray state transition is observed \citep[causing a decrease in flux;][]{veledina2013}. This is consistent with the time between the dip and the soft state transition observed for \source\ ($\simeq$1--2~weeks). Moreover, the hot flow model predicts that the optical/nIR spectrum hardens (i.e., becomes bluer) in this transition, consistent with the drop in $J-K$ color that we observed before the state transition (Figure~\ref{fig:lc}). \citet{veledina2011} interpret the optical/nIR secondary maxima seen after hard state transitions as the recovery of a hot inner flow \citep[see also][]{kalemci2013,veledina2013}. The same mechanism may account for the UV peak (and corresponding X-ray/UV anti-correlation) discussed in this work. However, in the picture of a recovering hot flow, the UV emission should rise before the nIR \citep[][]{veledina2013}, whereas in \source\ there are indications that the nIR started to rise first. Moreover, a delay as long as $\simeq$30 days between the state transition and the UV peak may be difficult to account for \citep[see also][]{kalemci2013}. It is worth noting that although UV monitoring is often hampered by substantial interstellar extinction, anti-correlated UV and X-ray emission has also been observed for the neutron star Cyg X-2 \citep[][]{rykoff2010}. On the other hand, such an anti-correlation was not seen for the black hole LMXBs \xte\ and \bhmontse, despite dense UV/X-ray coverage over a large range of fluxes \citep[][]{rykoff2007,armas2012}. The latter source may have never entered a soft state (\cite{armas2012}, see also \cite{corralsantana2013} and \cite{shahbaz2013}), which could possibly explain the absence of an anti-correlation given that the UV rise appears to be connected to a state transition. However, \xte\ was monitored for well over a month after it transitioned from the soft to the hard state, but its X-ray and UV emission remained tightly correlated \citep[][]{rykoff2007}. Future X-ray/UV monitoring efforts of suitable targets are warranted to further understand the role of the UV emission in state transitions. | 14 | 3 | 1403.0939 |
1403 | 1403.7034_arXiv.txt | We report our findings on a new quasi-periodic oscillation (QPO) and a long period from the ultraluminous X-ray source (ULX) X-2 in nearby galaxy NGC 4736 based on the {\it Chandra} and {\it XMM-Newton} archival data. To examine the timing properties, power density spectra of the source have been obtained using Fast Fourier Transform. Also the spectral parameters of the source have been calculated by obtaining and fitting the energy spectra. Power density spectrum of this source reveals a QPO peak at $0.73_{-0.14}^{+0.16}$ mHz with an fractional rms variability of $16\%$ using the {\it Chandra} data (in the year 2000-lower state of the source). The {\it XMM-Newton} data analysis indicates a peak at $0.53_{-0.35}^{+0.09}$ mHz with a fractional rms variation of $5\%$ (in the year 2006-higher state of the source). These recovered QPOs overlap within errors and may be the same oscillation. In addition, we detect a long periodicity or a QPO in the {\it Chandra} data of about $(5.2\pm2.0)\times10^{-5}$ Hz ($\sim$ 5.4 hrs) over 3 $\sigma$ confidence level. If this is a QPO, it is the lowest QPO detected from a ULX. The mass of the compact object in ULX X-2 is estimated using the Eddington luminosity and a disk blackbody model in the range (10$-$80) M$_{\sun}$. | Ultraluminous X-ray sources (ULXs) are off-nuclear point-like sources that have X-ray luminosities ($L_{x} \textgreater 3 \times 10^{39}$ erg $s^{-1}$) above the Eddington limit of a 20$M_{\sun}$ black hole (BH) (Feng \& Soria 2011). Although the true nature of these objects is not clear, current models propose several alternatives to explain their high luminosities: It could either be due to the geometric beaming if ULXs are powered by an accreting stellar black hole (King et al. 2001) or the super-Eddington fluxes originating from disks (Begelman 2002). Additionally, some ULXs could be powered by intermediate-mass black holes (Miller \& Colbert 2004). X-ray spectral and temporal analyses allow us to study the nature of ULXs and to constrain BH masses in ULXs (Soria \& Ghosh 2009). Quasi-Periodic Oscillations (QPOs) in X-ray binaries provide information about the inner accretion disk structure around the compact object (Mucciarelli et al. 2006). Up to now, QPOs have been detected from a few ULXs in nearby galaxies. The first ULX (X-1) that showed strong evidence for a QPO was in M82 (Strohmayer \& Mushotzky 2003). Using a 30 ks {\it XMM-Newton} observation they found that the power-density spectrum (PDS) of this source shows a prominent QPO peak at a frequency of 54 mHz with an integrated rms of 8.5\% in the (2$-$10) keV. They also determined a 107 mHz QPO from the same ULX using archived {\it RXTE} observations. Based on the Schwarzschild geometry, they calculated the mass of the BH in source X-1 as \textless 1.87$\times10^{4}$ $M_{\sun}$ using the highest QPO frequency. In the other studies using the {\it XMM-Newton} and {\it RXTE} observations similar properties were detected for the ULX in M82 (Mucciarelli et al. 2006, Feng \& Kaaret 2007). Also, Feng et al. (2010) reported QPOs at (3$-$4) mHz from a transient ULX (X42.3+59) in M82 using three {\it Chandra} and two {\it XMM-Newton} observations. They estimated the BH mass in the range of (1.2$-$4.3$)\times10^{4}$ $M_{\sun}$ in the ULX by scaling the QPO frequency to that of their type (A/B) of QPOs in stellar mass BHs. Liu et al. (2005) have presented the PDS of a ULX in NGC 628 (M74) showing a broad peak at a frequency range of (0.1$-$0.4) mHz in one {\it XMM-Newton} and two {\it Chandra} observations with an rms variation of 13.8\% and 23.9\% in (2$-$4) keV and (4$-$10) keV band, respectively. They estimated the mass of the compact object as $\sim$(2$-$20)$\times10^{3}$ $M_{\sun}$ using the scaling relation between break frequency and BH masses. Another ULX (Holmberg IX X-1) which showed a QPO was detected in Holmberg IX, using a 119 ks {\it XMM-Newton} observation (Dewangan et al. 2006b). It has a centroid frequency of 202.5 mHz with an rms of 6\% in the (0.2$-$10) keV energy band. However, Heil et al. (2009) reported that while the source Holmberg IX X-1 shows variability, it does not show a significant QPO feature. Strohmayer et al. (2007) found a ULX in NGC 5408 (X-1) which shows a pair of QPOs with a 4:3 ratio. The first peak is at 20 mHz and the second peak is at 15 mHz, at integrated rms of 9\% from a 130 ks {\it XMM-Newton} observation. Considering this QPO is analogous to the ones in Galactic systems (for example GRO J1655-40), they found a BH mass range of (1.5$-$3.5)$\times10^{3}$ $M_{\sun}$. Furthermore, using the mass-disk temperature scaling, $kT_{disk} \propto M^{-1/4}$, they also derived another mass range (1.81$-$4.74)$\times10^{3}$ $M_{\sun}$. Pasham and Strohmayer (2012) also presented detailed study of NGC 5408 X-1 using recent {\it XMM-Newton} observations. They detected QPOs in the range of (10$-$40) mHz with new observations. They also calculated a lower limit of $\sim800M_{\sun}$ on the mass of the BH in X-1 by scaling the minimum QPO frequency to a transition frequency of a reference stellar BH with known mass. Another ULX in nearby galaxy NGC 6946 shows possible QPO feature with a central frequency of $\sim$8.5 mHz (Rao et al. 2010). They calculated an integrated rms amplitude of 59\% for X-1 in the (1$-$10) keV energy range and a BH mass of $\sim 10^{3}$ $M_{\sun}$ for the compact source by scaling the frequency with mass. In the present work, we search the X-ray timing variations and spectral properties of the ULX X-2 in the LINER galaxy NGC 4736. The source is located in the disk of the galaxy. Akyuz et al. (2013) examined the long term light curves and energy spectrum of this source using {\it XMM-Newton} archival data revealing its transient nature. Previously, this source was detected as a point source using the {\it Chandra} archival data (Liu 2011) with no reference to its transient nature. A more recent study by Lin et al. (2013), also discusses the nature of the ULX X-2 using the {\it ROSAT}, {\it XMM-Newton} and {\it Chandra} archival data. They show highly variable dipping behaviour of the light curve in the brightest observations and discuss spectral properties in the dipping periods. They also examine the {\it HST} images and identify a point-like red optical counterpart candidate. Considering the colors and the luminosities of this candidate, they suggest that it could be a G8 supergiant or a dwarf star in the globular cluster. Differently from the Lin et al. (2013) work, we concentrate on the timing properties and search for periodicities from the ULX X-2. This paper organized as follows: The observations and methods used in data reductions are described in Sect. 2. The timing and spectral analysis are given in Sect. 3 and Sect. 4. Discussions on the QPOs and the BH mass estimations are given in Sect. 5. \begin{figure}[t] \includegraphics[scale=0.46, angle=0]{avdan_fig1.ps} \caption{The X-ray light curves of X-2 in the (0.2$-$10) keV energy band. Observation dates are labelled on the light curves. The dashed lines show the mean count rate. High background flare times are excluded from the ObsID 0404980101. Both light curves have a bin size of 1000 s.} \end{figure} | In this work, we have presented the energy and power spectra of the ULX X-2 in NGC 4736 using archival {\it Chandra} and {\it XMM-Newton} observations. {\it Chandra} observation of the source revealed a QPO with a centroid frequency of $\nu_{QPO}=0.73_{-0.14}^{+0.16}$ mHz (with an integrated rms fractional variability of $16\% \pm 3\%$ and $Q = 7.3$) while the {\it XMM-Newton} observation showed a QPO at $0.53_{-0.35}^{+0.09}$ mHz (with an integrated rms of $5\% \pm 1\%$ and a $Q = 5.3$). The two frequencies overlap within their error limits, thus these could be either slightly differing QPOs or the same QPO from the system. It is consistent that the low frequency QPO has diminished in rms percent variability when the source changed to a higher luminosity state with the outer parts of the disk being more stable to variations. It is expected that in high states of BH LMXBs low frequency QPOs are suppressed. The energy spectrum for the {\it Chandra} observation shows a best fit with a power-law model (${\Gamma}$ $\sim 2.5$). The {\it XMM-Newton} spectrum is best fitted with a two-component model, the disk blackbody plus power-law ($T_{in}$=0.7 keV, ${\Gamma}$ $\sim$ 1.7) as published earlier by Akyuz et al. (2013). Our analysis showed that the source luminosity has changed by a factor $\sim$20 times between two observations. There are also couple of transient ULXs that exhibits decrease in luminosity below the ULX regime: CXOM31 J004253.1+411422 in M31 (Kaur et al. 2012), ULX in M83 (Soria et al. 2012), 1RXH J132519.8-430312 and CXOU J132518.2-430304 in NGC 5128 (Burke et al. 2013). The level of change in luminosity of the source is not well-known since the source was observed only four times within eight years. Its spectral shape has changed from a one-component model of a PL to a two-component model of PL+DISKBB where the disk component contributed $\sim60\%$ and the power-law component yielded the rest of the total flux. This indicates that X-2 may be in a spectral state analogous to the thermal state of Galactic BH X-ray binaries (Remillard \& McClintock 2006) during the {\it XMM-Newton} observation. The source has a steep PL photon index as calculated from the {\it Chandra} observation with a lower luminosity than the thermal state. However, the steep power law state of Galactic BH X-ray binaries is characterized by a relatively high luminosity with a PL component of ${\Gamma} > 2.4$ (Remillard \& McClintock 2006). The ULX X-2 exhibits opposite behaviour, however a similar unusual low/soft state has also been observed from ULX X-2 in NGC 1313 (Feng \& Kaaret 2006). This unusual behaviour may arise from the lack of good statistics in the {\it Chandra} data. On the other hand, Lin et al. (2013) also carried out spectral analysis of X-2 using one {\it ROSAT}, one {\it XMM-Newton} and two {\it Chandra} observations. Their best-fitting spectral models, PL model in the {\it Chandra} (ObsID 808) and PL+MCD model in the {\it XMM-Newton} (ObsID 0404980101) data, and the derived spectral parameters are in agreement with ours. They noticed that the source might be in the hard state instead of the steep power-law, if one takes into account the large uncertainty of PL photon index ($\sim 2.5 \pm0.5$) due to the poor quality of the data. The mass of the BH in ULX systems can be calculated using the relation M=$c^{3}/(2\pi 6^{3/2}G {\nu}_{QPO})$ $\simeq$ 2190/${\nu}_{QPO}$ $M_{\sun}$ (van der Klis 2006). This is based on the assumption that the QPO frequencies are associated with the Keplerian frequency at the innermost circular orbit around a Schwarzschild BH. The formula yields a mass value of $\sim$ $3.5\times 10^{6}$ $M_{\sun}$ for X-2, which is rather high for even an intermediate BH. Another approach is to consider the inverse proportionality between the BH mass and the QPO frequencies, and how this scales over different systems. Dewangan et al. (2006a) used this proportionality for ULX X-1 in M82 and found the mass of the BH to be in the range of (25$-$520)$M_{\sun}$. We used the same scaling argument to approximate the mass of the BH in X-2 of NGC 4736. We chose the ULX in NGC 628 for the scaling since the detected QPOs have similar low frequencies and luminosities ($L_{x}=(4.5-13.4)\times10^{38}$ erg s$^{-1}$, Liu et al. 2005). Also, this source is one of the ULX which have the lowest QPO frequency ever detected. Using this scaling argument and assuming the mass ($(2-20)\times 10^{3}$ $M_{\sun}$) and QPO frequencies ($(1-4)\times 10^{-4}$ Hz) of the ULX in NGC 628 (Liu et al. 2005), we estimated a BH mass of $(2-400)\times 10^{2}$ $M_{\sun}$ for the compact source in X-2. However, we note that some ULXs have low luminosities and they were estimated to have stellar/massive-stellar ($10 M_{\sun} \leq M \leq 100M_{\sun}$) BHs (Kaur et al. 2012; Soria et al. 2012). On the other hand, the possible mass for the compact object in X-2 can be estimated assuming the source emits at the Eddington limit. The BH mass of $M_{bh}$ $\sim$ 10 $M_{\sun}$ is found by using the highest luminosity value ($\sim$1.7$\times$10$^{39}$ erg s$^{-1}$) with this assumption. Also considering the dominant contribution from the disk and the disk component model parameter, the inner disk radius can be derived as $R_{in} \sim$ 96$(\cos \theta)^{-0.5}$ km. Then, an upper limit for the compact source mass is estimated as $\leq$ 80 $M_{\sun}$ in X-2 using the inner disk radius (Makishima et al. 2000). Considering this mass range of (10$-$80) $M_{\sun}$, the compact source in X-2 is probably a stellar mass BH ($M \leq 20 M_{\sun}$) or a massive-stellar BH ($20 M_{\sun} \leq M \leq 100M_{\sun}$)(Feng \& Soria 2011). In addition, we have detected a long periodicity of $(5.2\pm2.0)\times10^{-5}$ Hz ($\sim$5.4 hrs) above 3 $\sigma$ confidence level in the low state of the source. We speculate that this may be another low frequency QPO from the system that may possibly be consistent with an intermediate mass BH scenario owing to the plausible large size disk. In such a case, this may be the lowest QPO detected from a ULX. However, it is also possible that this is the orbital period of the underlying binary which, then, it will be more consistent with a stellar/massive-stellar size BH. We note that the source shows low and high states of luminosity and there is not enough data on the source to conclusively decide whether it belongs to XRB class or a ULX classification. We encourage further monitoring X-ray observations of X-2 in the nearby galaxy NGC 4736 to understand its true physical parameters and characteristics. In addition, these should be aided with observations in the other wavelength ranges (optical, IR, radio), as well. | 14 | 3 | 1403.7034 |
1403 | 1403.2757_arXiv.txt | We present a family of solutions of Einstein's gravity minimally coupled to a complex, massive scalar field, describing asymptotically flat, spinning black holes with scalar hair and a regular horizon. These hairy black holes (HBHs) are supported by rotation and have no static limit. Besides mass $M$ and angular momentum $J$, they carry a conserved, continuous Noether charge $Q$ measuring the scalar hair. HBHs branch off from the Kerr metric at the threshold of the superradiant instability and reduce to spinning boson stars in the limit of vanishing horizon area. They overlap with Kerr black holes for a set of $(M,J)$ values. A single Killing vector field preserves the solutions, tangent to the null geodesic generators of the event horizon. HBHs can exhibit sharp physical differences when compared to the Kerr solution, such as $J/M^2>1$, quadrupole moment larger than $J^2/M$ and larger orbital angular velocity at the innermost stable circular orbit. Families of HBHs connected to the Kerr geometry should exist in scalar (and other) models with more general self interactions. | 14 | 3 | 1403.2757 |
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1403 | 1403.0564_arXiv.txt | The modeling of turbulence, whether it be numerical or analytical, is a difficult challenge. Turbulence is amenable to analysis with linear theory if it is subject to rapid distortions, i.e., motions occurring on a time scale that is short compared to the time scale for non-linear interactions. Such an approach (referred to as rapid distortion theory) could prove useful for understanding aspects of astrophysical turbulence, which is often subject to rapid distortions, such as supernova explosions or the free-fall associated with gravitational instability. As a proof of principle, a particularly simple problem is considered here: the evolution of vorticity due to a planar rarefaction in an ideal gas. Analytical solutions are obtained for incompressive modes having a wave vector perpendicular to the distortion; as in the case of gradient-driven instabilities, these are the modes that couple most strongly to the mean flow. Vorticity can either grow or decay in the wake of a rarefaction front, and there are two competing effects that determine which outcome occurs: entropy fluctuations couple to the mean pressure gradient to produce vorticity via baroclinic effects, whereas vorticity is damped due to the conservation of angular momentum as the fluid expands. Whether vorticity grows or decays depends upon the ratio of entropic to vortical fluctuations at the location of the front; growth occurs if this ratio is of order unity or larger. In the limit of purely entropic fluctuations in the ambient fluid, a strong rarefaction generates vorticity with a turbulent Mach number on the order of the root-mean square of the ambient entropy fluctuations. The analytical results are shown to compare well with results from two- and three-dimensional numerical simulations. Analytical solutions are also derived in the linear regime of Reynolds-averaged turbulence models. This highlights an inconsistency in standard turbulence models that prevents them from accurately capturing the physics of rarefaction-turbulence interaction. In addition to providing physical insight, the solutions derived here can be used to verify algorithms of both the Reynolds-averaged and direct numerical simulation variety. Finally, dimensional analysis of the equations indicates that rapid distortion of turbulence can give rise to two distinct regimes in the turbulent spectrum: a distortion range at large scales where linear distortion effects dominate, and an inertial range at small scales where non-linear effects dominate. | Introduction} Subsonic turbulence is important in both the intracluster and intergalactic medium \citep{sch04,sub06}. Supersonic turbulence is present in the interstellar medium and plays an indispensable role in star formation \citep{mo07}. Numerical modeling of turbulence is difficult and fraught with uncertainty \citep{bs12}, and any analytical results that can be obtained provide both a check for numerical codes and a wider view of parameter space (restricted by the assumptions underlying the analytical results). Although analytical modeling comes with its own set of difficulties, significant progress can be made for turbulence subjected to rapid distortions. Rapid distortion theory (RDT) is an analytical approach to the study of turbulence for conditions under which non-linear effects can be neglected \citep{sav87}. Such conditions pertain, for example, to a supernova explosion propagating through a turbulent medium or to turbulent eddies in gravitational free-fall. The purpose of this work is to investigate a particularly simple problem using RDT as a proof-of-principle for its application to more realistic astrophysical flows. The problem to be studied is the evolution of subsonic turbulence in an ideal gas subject to a centered rarefaction. Such a flow occurs, for example, when a shock propagates from heavy to light material in an interaction with a contact discontinuity \citep{mik94}. To make the problem analytically tractable, the analysis is restricted to modes that are oriented perpendicular to the distortion; these are the incompressive modes that couple most strongly to the mean flow. A complete RDT analysis of rarefaction-turbulence interaction would need to take into account the full spectrum of linear modes. Despite this restriction, the one-dimensional analytical solution derived here captures the essential physics. Favorable comparisons are made with both two- and three-dimensional numerical simulations. Another approach to modeling turbulence when sufficient resolution is not available is to employ a Reynolds-averaged turbulence model \citep{gb90,dt06,ms13,ms14}. These models have been used, for example, to capture mixing in interactions between active galactic nuclei and bubbles \citep{sb08}, between high-redshift galactic outflows and clouds \citep{gs11}, between shocks and clouds \citep{pit09,pit10}, and between galactic haloes and the intergalactic medium \citep{clo13}. Turbulence modeling comes with its own set of uncertainties, due to multiple closures and poorly-constrained model coefficients. Analytical solutions are derived here for Reynolds-averaged models in the linear regime which can serve as a verification test for these models. In addition, comparison to the analytical linear theory highlights an inconsistency in standard models that prevents them from correctly capturing the physics of rarefaction-turbulence interaction. A simple proposal for correcting this inconsistency will be provided. \S\ref{BEMF} outlines the basic equations and provides the well-known expressions for a centered rarefaction. An overview of RDT along with its application to the problem at hand is given in \S\ref{PCE}. Comparisons between RDT and numerical simulations are provided in \S\ref{NR}, Reynolds-averaged models are discussed in \S\ref{RANS}, and \S\ref{SD} summarizes the analysis and gives suggestions for future work. | 14 | 3 | 1403.0564 |
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1403 | 1403.2561_arXiv.txt | {Galaxy Evolution Explorer (GALEX), the first all sky imaging ultraviolet (UV) satellite, has imaged a large part of the sky providing an excellent opportunity for studying UV star counts. Combining photometry from the different wavelengths in the infrared (from Wide-field Infrared Survey (WISE) and Two Micron All Sky Survey (2MASS)) to UV allows us to extract a real star catalog from the GALEX source catalog.} {The aim of our study is to investigate in detail the observed UV star counts obtained by GALEX vis-a-vis the model simulated catalogs produced by the Besan\c{c}on model of stellar population synthesis in various Galactic directions, and to explore the potential for studying the structure of our Galaxy from images in multiple near-UV (NUV) and far-UV (FUV) filters of the forthcoming Ultraviolet Imaging Telescope (UVIT) to be flown onboard ASTROSAT.} {We have upgraded the Besan\c{c}on model of stellar population synthesis to include the UV bands of GALEX and UVIT. Depending on the availability of contiguous GALEX, Sloan Digital Sky Survey (SDSS), WISE and 2MASS overlapping regions, we have chosen a set of nineteen GALEX fields which spread over a range of Galactic directions. We selected a sample of objects from the GALEX database using the \emph{CASjobs} interface and then cross-matched them with the WISE+2MASS and SDSS catalogs. UV stars in the GALEX catalog are identified by choosing a suitable infrared (IR) colour, $J - W1$ (W1 is a WISE band at 3.4 $\mu$m), which corresponds to a temperature range from 1650 K to 65000 K. The IR colour cut method, which is used for the first time for separation of stars, is discussed in comparison with the GALEX+SDSS star counts method.} {We present the results of the UV star counts analysis carried out using the data from GALEX. We find that the Besan\c{c}on model simulations represent the observed star counts of both the GALEX All-sky Imaging Survey (AIS) and Medium Imaging Survey (MIS) well within the error bars in various Galactic directions. Based on the model analysis, we separated out white dwarfs (WDs) of the disc and blue horizontal branch stars (BHBs) of the halo from the observed sample by selecting a suitable $FUV - NUV$ colour.} {The Besan\c{c}on model is now ready for further comparisons in the UV domain and will be used for prospective studies for the UVIT instrument to be flown onboard ASTROSAT.} | The Milky Way is the best studied Galaxy in the universe; its structure and evolution have been studied by a variety of techniques. In the early 20$^{th}$ century, \citet{Kapteyn22} first studied the geometrical structure of the Galaxy using the star counts method whereby he counted stars on the photographic plates in selected areas of the sky. Since then the star counts method has been used as one of the preferred methods to constrain the structural parameters of the Galaxy effectively. Several reviews \citep{Bahcall86,Freeman87,Gilmore88,Majewski93,Helmi08,Ivezic12} have discussed the connection of star counts to the Galactic structure. The advent of instruments with better resolution and greater sensitivity have enabled us to obtain photometric observations covering large parts of the sky in several wavelength bands. The population synthesis models of the Milky Way are well supported by these observations in predicting the different structural parameters of the Galaxy, such as stellar densities, scale length, scale height, etc. Among the models built to understand the Galactic structure by star counting method, one can cite: \citet{Bahcall80}, \citet{Gilmore83}, \citet{Robin86}, \citet{Robin03}, \citet{Girardi05} and \citet{Juric08}. However, the above Galaxy models are predominantly based on the visible and IR photometric surveys. Very few attempts \citep{Brosch91,Cohen94} had been made to study the star counts in UV prior to GALEX due to a lack of availability of UV photometric surveys. The advent of GALEX, which provided a wide sky coverage in UV, now allows new analysis of the UV sky \citep[][ among others]{Xu05,Bianchi11a,Bianchi11b,Bianchi13}. An attempt has also been made to predict the star counts in the X-ray band \citep{Guillout96} by extending the Besan\c{c}on model of stellar population synthesis \citep{Robin86} to the ROSAT PSPC energy bands. Indeed, the UV surveys, among others, could help in tracing the spiral structures which mainly contain very young stars. The UV surveys also help in constraining the shape of the initial mass function (IMF) towards the high-mass star end as well as elucidating the recent star formation history. Moreover, they also trace very blue populations such as WDs and BHBs deep in the halo population, which in turn trace the streams and relics of ancient accretion in the Milky Way halo. GALEX has covered a large part of the sky which provides an opportunity to explore and characterize these hot sources in the FUV (1344 - 1786 \AA, $\lambda_\mathrm{eff}$ = 1538.6 \AA) and NUV (1771 - 2831 \AA, $\lambda_\mathrm{eff}$ = 2315.7 \AA) wavebands with better resolution and greater sensitivity than the previous surveys. A vivid description of the source selection, FUV and NUV magnitude error cuts and the statistical analysis of the GALEX catalog is provided by \citet{Bianchi07}, \citet{Bianchi09} and \citet{Bianchi11a, Bianchi13}. Detection of WDs and BHBs is one of the main achievements of GALEX as these sources are elusive in the other wavelength bands of the electromagnetic spectrum due to their high temperature. WDs and BHBs are integral to the study of stellar evolution and structure of the Milky Way as they belong to different stellar populations of the Galaxy. We have upgraded the Besan\c{c}on model of stellar population synthesis to include the UV bands of GALEX and the upcoming UVIT\footnote{http://www.iiap.res.in/Uvit} (which will be flown onboard ASTROSAT) to predict star counts in different parts of the sky \citep{Todmal10}. UVIT will image the sky in the FUV (1300 - 1800 \AA) and NUV (2000 - 3000 \AA) channels, each having five filters, at a high resolution of \(1.8''\) \citep{Postma11,Kumar12a,Kumar12b}. Better positional accuracy of UVIT as compared to GALEX will enable more reliable cross correlation with other catalogs which will be of great utility in inferring the Galactic structure using the star counts technique. The transmission curves (effective area versus wavelength) for the FUV and NUV bands of GALEX together with each of the five FUV (left panel) and NUV (right panel) filters of the upcoming UVIT/ASTROSAT are shown in Figure 1. We have included the effective area curves of both the GALEX and all the UVIT/ASTROSAT bands in the model to simulate the UV star counts in these bands. Apart from the GALEX bands, we will discuss the model simulated star counts of the BaF2 (FUV: 1370 - 1750 \AA, $\lambda_\mathrm{eff}$ = 1504 \AA) and NUVB4 (NUV: 2505 - 2780 \AA, $\lambda_\mathrm{eff}$ = 2612 \AA) bands of UVIT/ASTROSAT. The expected sensitivity limits (5$\sigma$) in AB magnitude system in the UVIT BaF2 (FUV) and NUVB4 (NUV) wavebands, for an exposure time of 200 seconds, are 20.0 and 21.2 magnitudes, respectively (ASTROSAT Handbook 2013; private communication). It is worth mentioning here that throughout the paper we have used AB system for the GALEX, UVIT and SDSS data sets, whereas the 2MASS and WISE data sets are in the Johnson system (see Section 2). We give details of the observations and selection of UV stars in Section 2. We describe about the Besan\c{c}on model Galaxy model in Section 3 and discuss the comparison of the GALEX+WISE+2MASS and GALEX+SDSS star counts in Section 4. We present the comparison of the model with the observations in Section 5, and discuss the distribution of the model star counts in Section 6. We mention the identification of WDs and BHBs using $FUV - NUV$ colour in Section 7. Finally, we summarize our conclusions in Section 8. | The Besan\c{c}on model of stellar population synthesis has been previously checked at many different wavelengths from visible (U band) to mid-IR (12 $\mu$m). The model produces accurate star counts up to magnitude $\sim$ 22 in the visible or 18 in the K band. However, the stars that dominate the counts in the UV were not previously checked vis-a-vis model predictions. The availability of the GALEX data gives opportunity to check model predictions for high temperature, blue stars, specially BHBs from the halo and WDs from the disc. We have shown that the model performs very well for these types of stars as it does for other types. The model provides a good check that the population synthesis scheme gives predictions which are consistent with each other at all wavelengths. To do so, we make use of \citet{Holberg06} models which provide good stellar atmospheres and cooling tracks for WDs. However, the ratio between DA and DB type WDs has to be investigated more deeply. We have generally considered a simple dust distribution while limiting the comparisons to $|b| > 20\degr$. In future, we will compare the model at lower latitudes, in particular for the sake of analysis of the spiral structure, assuming the 3D extinction map from \citet{Marshall06}. We also compared predictions in the UV bands from the TRILEGAL model with our model and found that the predictions of the Besan\c{c}on model are in better agreement with the observation than the TRILEGAL model as shown in Figure 10. However, in the faintest NUV magnitude bins TRILEGAL seems to be better in the GAR field. It will be something to look at carefully in the future using the all sky observations of GALEX and WISE, and we aim to present a detailed comparison between observations and the model. We plan to complete the analysis by comparing model predictions with a variety of models of WDs, varying the tracks and investigating whether it could be possible to constrain the star formation history of the disc from the WDs distribution. Moreover, an analysis of the thick disc WD luminosity function could also be interesting for constraining the formation history of this old population, but it would require complementary kinematical data. We have seen that BHBs are a major component of GALEX stars. An analysis of this component could lead to constraints on the shape of the halo, once the contamination by extra-galactic objects is eliminated. The final model can be safely used to predict star counts of various types in the UV wavelengths at the level of a few percent in many Galactic directions; the model produces star counts that match well down to FUV $\sim$ 20.0, NUV $\sim$ 20.5 for AIS, and FUV $\sim$ 22.5, NUV $\sim$ 22.0 for MIS. However, for the hot WDs, there is a mismatch of UV colours between the model and observation. A more detailed study is planned to explain the discrepancies by changing the WD luminosity function and the scale lengths alternatively. A study is also going on to better constrain the thick disc shape from large surveys in the visible and near-IR (Robin et al., in prep). We plan to further investigate the UV star counts with this revised model and the GALEX survey in the near future. The Besan\c{c}on model is also developed to predict star counts in the UV passbands of the forthcoming UVIT telescope to be flown onboard ASTROSAT. We compared the model-predicted star counts at two of the UVIT filters with that of the GALEX observed star counts because of the similar wavelength coverage of both the instruments. The UVIT-predicted star counts are sensitively different from the GALEX observed star counts due to the differences in effective wavelengths. UVIT star counts will be very useful to separate out different stellar populations since they have several UV colours and better angular resolution compared to GALEX, which in turn will help us to estimate the structural parameters of the Galaxy with better precision. | 14 | 3 | 1403.2561 |
1403 | 1403.3617_arXiv.txt | We have obtained extensive high-quality spectroscopic observations of the OGLE-LMC-CEP-1718 eclipsing binary system in the Large Magellanic Cloud which Soszynski et al. (2008) had identified as a candidate system for containing two classical Cepheids in orbit. Our spectroscopic data clearly demonstrate binary motion of the Cepheids in a 413-day eccentric orbit, rendering this eclipsing binary system the first ever known to consist of two classical Cepheid variables. After disentangling the four different radial velocity variations in the system we present the orbital solution and the individual pulsational radial velocity curves of the Cepheids. We show that both Cepheids are extremely likely to be first overtone pulsators and determine their respective dynamical masses, which turn out to be equal to within 1.5 \%. Since the secondary eclipse is not observed in the orbital light curve we cannot derive the individual radii of the Cepheids, but the sum of their radii derived from the photometry is consistent with overtone pulsation for both variables. The existence of two equal-mass Cepheids in a binary system having different pulsation periods (1.96 and 2.48 days, respectively) may pose an interesting challenge to stellar evolution and pulsation theories, and a more detailed study of this system using additional datasets should yield deeper insight about the physics of stellar evolution of Cepheid variables. Future analysis of the system using additional near-infrared photometry might also lead to a better understanding of the systematic uncertainties in current Baade-Wesselink techniques of distance determinations to Cepheid variables. | Classical Cepheids are distance indicators par excellence and a fundamental rung on the cosmic distance ladder, connecting our Milky Way galaxy to galaxies in the Local Group and beyond (Freedman et al. 2001; Gieren et al. 2005a, 2006; Pietrzynski et al. 2006; Riess et al. 2011). In order to render Cepheids even more robust and reliable distance indicators, it is imperative to understand their physical and evolutionary properties with the highest possible accuracy. In that context, it has been a breakthrough to find classical Cepheids in detached, double-lined eclipsing binary systems which permit a determinion of their basic physical parameters much more accurately than what is possible for any single Cepheid star. In particular, the analysis of the OGLE-LMC-CEP-0227 system located in the Large Magellanic Cloud (LMC), containing a classical Cepheid pulsating with a period of 3.8 days together with a stable red giant in a 310-day orbit, has yielded for the first time a Cepheid mass and radius determination accurate to 1 \%, and valuable independent insight on the p-factor needed for Baade-Wesselink-type analyses (Pietrzynski et al. 2010; Pilecki et al. 2013). A second eclipsing binary system in the LMC containing an even shorter-period classical Cepheid, OGLE-LMC-CEP-1812, was analyzed by Pietrzynski et al. (2011) and again yielded a very accurate measurement of the dynamical mass of the Cepheid. These two Cepheid mass determinations have gone a long way to solve the famous Cepheid mass discrepancy problem, leading to improved predictions of Cepheid masses from stellar pulsation and evolution theories. (Marconi et al. 2013; Prada Moroni et al. 2012). In the present paper, we report on the confirmation of an even more exotic, and so far unique, eclipsing binary system in the LMC consisting of a {\it pair of classical Cepheids in a 413-day orbit}. The system, herein named OGLE-LMC-CEP-1718, was discovered and identified as a {\it double Cepheid} by Alcock et al. (1995). Later Soszynski et al. (2008) found that it also exhibits eclipsing variability, but it was not yet clear if the two Cepheids were indeed gravitationally bound. Our spectroscopic observations of this double-lined system over the past years clearly show that the two Cepheids orbit each other, with the additional radial velocity variability of the Cepheids due to their pulsations superimposed on their orbital radial velocity curves. Evidently, the analysis of this system and the characterization of the physical properties of its coeval Cepheids holds great promise to deepen our understanding of Cepheid physics and evolution. In this paper, we present spectroscopic observations of OGLE-LMC-CEP-1718 and extract the orbital radial velocity curves of the two components of the system and the individual pulsational radial velocity curves of the Cepheids. We add new photometric data from OGLE III and OGLE IV surveys to that presented by Soszynski et al. (2008) to the radial velocity data to obtain the orbital solution as well as a determination of several physical parameters of the Cepheids, particularly their masses and pulsation modes. | From our orbital solution, we find that the masses of the Cepheids are 3.3 and 3.28 $M_\odot$, respectively, individually determined with an accuracy of 3 \% (see Table 3). Because the high eccentricity error does not contribute to the evaluation of the mass ratio, it is determined with a much better accuracy. The analysis indicates that the two Cepheids in the OGLE-LMC-CEP-1718 system have equal masses to within $\pm$ 1.5 percent. The very short pulsation periods suggest pulsation in non-fundamental modes. In Figure 7 we have plotted the positions of the two Cepheids on the I-band light curve Fourier decomposition diagrams of Soszynski et al. (2008). The loci of both stars on these diagrams, particularly on the $R_{21}-\log P$ diagram, strongly suggest that both Cepheids in OGLE-LMC-CEP-1718 are pulsating in the first overtone mode. This result can be checked in a different way. Since we cannot observe the secondary eclipse in the light curve, we cannot determine the individual radii of the Cepheids. However the analysis of the light curve does return the sum of the radii, in this case 52.5 $\pm$ 1.5 $R_{\odot}$. Assuming first overtone pulsation for the two Cepheids in OGLE-LMC-CEP-1718, we can calculate their expected radii from a period-radius relation calibrated for first overtone Cepheids. Using the observational relation given by Sachkov (2002), we obtain radii of R = (26.1 $\pm$ 2.5) $R_{\odot}$ for the primary, and R = (31.0 $\pm$ 2.6) $R_{\odot}$ for the secondary (longer-period) Cepheid, with the radii ratio of $1.19$. These predictions are in excellent agreement with the predictions from the theoretical period-radius relation for first overtone Galactic Cepheids of Bono et al. (2001; 27 and 32 $R_{\odot}$, respectively). As the ratio is much better constrained than the radii themselves we have used it to calculate the individual radii of the Cepheids (using the known sum), obtaining $R_1=24 R_{\odot}$ and $R_2=28.5 R_{\odot}$ which is clearly consistent with the values from the given relation. If we assume fundamental mode pulsation for both Cepheids and use the fundamental mode Cepheid period-radius relation from Sachkov (2002) (which is very similar to other calibrations of that relation, e.g. Gieren et al. 1998), the expected sum of the radii is 44.0 $\pm$ 0.3 $R_{\odot}$, indicating that fundamental mode pulsation is much more unlikely. We have to note however, that in this case the calculations are based on the extrapolation as the periods of our stars are shorter than the shortest one among the stars used to obtain the relation. Yet another argument supporting the first overtone pulsation hypothesis comes from the observed brightness of OGLE-LMC-CEP-1718. Using the fitted PL relations for first overtone Cepheids in the LMC from the OGLE project (Soszynski et al. 2008), the expected apparent magnitudes for the primary Cepheid are 15.439 and 16.110 in I and V bands, respectively, whereas for the secondary Cepheid the corresponding values are 15.104 and 15.786. This leads to expected total apparent magnitudes of both components of $I_{tot}$ = 14.506 and $V_{tot}$ = 15.183, respectively. The observed apparent magnitudes of the system are $I_{obs}$ = 14.511 and $V_{obs}$ = 15.190, in excellent agreement with the expected magnitudes if both Cepheids pulsate in the first overtone mode. One possibility to reconcile the fact that both Cepheids in the system have the same masses, but different pulsation periods, would be the assumption that the primary, shorter-period Cepheid is actually pulsating in the second overtone mode. The observed period ratio of 1.96/2.48 = 0.79 would be consistent with the hypothesis that the primary Cepheid is pulsating in the second overtone while the secondary Cepheid is a first-overtone pulsator-a period ratio of about 0.8 is indeed commonly observed for double-mode 1O/2O Cepheids. However, in the OGLE database of Cepheids in the Magellanic Clouds (Soszynski et al. 2008) which contains the largest samples of Cepheids, among others about 100 single-mode second-overtone Cepheids and about 420 double-mode 1O/2O Cepheids, the largest known period of a second overtone Cepheid is 1.32 days (in the double-mode Cepheid OGLE-SMC-CEP-0305 in the Small Magellanic Cloud). This is very much shorter than the period of 1.96 days observed for the primary Cepheid in our OGLE-LMC-CEP-1718 system. The largest known amplitude of the I-band light curve of a second overtone Cepheid is 0.138 mag (in the single-mode Cepheid OGLE-SMC-CEP-3509). Our object has a smaller amplitude of 0.097 mag (see Fig. 4), but the amplitude is decreased by the light from the secondary component through blending. Transforming the magnitudes to fluxes and removing the contribution of the secondary Cepheid, and transforming the flux back to magnitudes now yields an I-band amplitude of 0.231 mag for the short-period primary Cepheid in the OGLE-LMC-CEP-1718 system. This is much larger than any known amplitude of a second-overtone oscillation. The shapes of the light curves cannot be directly compared because there are no second-overtone Cepheids with periods around 2 days. However, it can be stated that all second-overtone Cepheid light curves are more symmetrical than the one of the primary in our system, they are indeed nearly sinusoidal. Finally, the total luminosity of the OGLE-LMC-CEP-1718 system perfectly agrees with the assumption that the system consists of two first-overtone Cepheids, as already mentioned above. In conclusion, if the 1.96-day Cepheid in our system would indeed be a second overtone pulsator, it would be the longest-period and the largest-amplitude second-overtone Cepheid known in any galaxy. Given the large number of second-overtone Cepheids known to-date, it seems extremely unlikely that we found such an extreme object in the only eclipsing binary system consisting of two classical Cepheids that has been discovered so far. Our conclusion then is that we have found a system composed of two classical Cepheids which have within a 1.5 \% uncertainty identical masses, both stars are almost certainly pulsating in the first overtone mode, presumably have the same ages, but have substantially different periods and luminosities. It will be challenging for stellar evolutionary theory to explain the observed properties of these Cepheids, which will be the topic of a forthcoming study of our group. | 14 | 3 | 1403.3617 |
1403 | 1403.5611_arXiv.txt | In this paper, we give a detailed user's guide to the \amidas\ (A Model--Independent Data Analysis System) package and website, which is developed for online simulations and data analyses for direct Dark Matter detection experiments and phenomenology. Recently, the whole \amidas\ package and website system has been upgraded to the second phase: \amidasii, for including the new developed Bayesian analysis technique. \amidas\ has the ability to do full Monte Carlo simulations as well as to analyze real/pseudo data sets either generated by another event generating programs or recorded in direct DM detection experiments. Moreover, the \amidasii\ package can include several ``user--defined'' functions into the main code: the (fitting) one--dimensional WIMP velocity distribution function, the nuclear form factors for spin--independent and spin--dependent cross sections, artificial/experimental background spectrum for both of simulation and data analysis procedures, as well as different distribution functions needed in Bayesian analyses. | Weakly Interacting Massive Particles (WIMPs) $\chi$ arising in several extensions of the Standard Model of electroweak interactions are one of the leading candidates for Dark Matter (DM). Currently, direct DM detection experiments based on measuring recoil energy deposited in a low--background underground detector by elastic scattering of ambient WIMPs off target nuclei are one of the most promising methods for understanding DM properties, identifying them among new particles produced (hopefully in the near future) at colliders, as well as studying the (sub)structure of our Galactic halo (for reviews, see Refs.~\cite{SUSYDM96, Bertone05, Bergstrom12}). Since 2007 we develop a series of new methods for analyzing data, i.e.~measured recoil energies, from % direct detection experiments as model--independently as possible. Up to now we could in principle reconstruct the (moments of the) one--dimensional velocity distribution function of halo WIMPs \cite{DMDDf1v, DMDDf1v-Bayesian}, as well as determine the WIMP mass \cite{DMDDmchi} and (ratios between) different WIMP couplings/cross sections on nucleons \cite{DMDDfp2, DMDDranap}. Following the development of these model--independent data analysis procedures, we combined the code for our simulations to a compact system: \amidas\ (A Model--Independent Data Analysis System). Meanwhile, we modified the code to be able to analyze external data sets either generated by another event generating programs or recorded in direct DM detection experiments. Under the collaboration with the ILIAS Project \cite{ILIAS} and the Dark Matter Network Exclusion Diagram (DAMNED, a.k.a.~Dark Matter Online Tools) \cite{DAMNED}, an online system has also been established in January 2009 \cite{AMIDAS-web, AMIDAS-web-TiResearch}, in order to offer an easier, more convenient and user--friendly environment for simulations and data analyses for (in)direct DM detection experiments and phenomenology. In the first phase of the \amidas\ package and website, the options for target nuclei, for the velocity distribution function of halo WIMPs, as well as for the elastic nuclear form factors for spin--independent (SI) and spin--dependent (SD) WIMP--nucleus interactions are fixed and only some commonly used analytic forms have been defined \cite{AMIDAS-SUSY09}. Users can not choose different detector materials nor use different WIMP velocity distribution function nor nuclear form factors for their simulations and/or data analyses. In order to offer more flexible use of the WIMP velocity distribution as well as the nuclear form factors, the \amidas\ package has been extended to be more user--oriented and able to include {\em user--uploaded} files for defining their own functional forms in September 2009 \cite{AMIDAS-f1vFQ}% \footnote{ Note that, since the \amidas\ code has been written in the C programming language, all user--defined functions to be included into the \amidas\ package must be given in the syntax of C. On the other hand, for drawing output plots, since the {\tt Gnuplot} package has been used in \amidas, the uploaded files for drawing e.g.~the predicted spectrum of measured recoil energy, must be given in the syntax of {\tt Gnuplot} \cite{gnuplot}. More detailed descriptions will be given in Secs.~3 and 5 as well as in Appendix. }. Besides of the purposed model--independent reconstructions of different WIMP properties \cite{DMDDf1v, DMDDf1v-Bayesian, DMDDmchi, DMDDfp2, DMDDranap}% \footnote{ Note that our model--independent methods are developed for reconstructing WIMP properties with {\em positive} signals (and probably small fractions of unrejected background events in the analyzed data sets \cite{DMDDbg-mchi, DMDDbg-f1v, DMDDbg-fp2, DMDDbg-ranap}). With however {\em negative} results or before enough (${\cal O}(50)$ to ${\cal O}(500)$) WIMP signals could be accumulated, other well--built model--independent methods developed in Refs.~% \cite{Fox10a, Fox10b, Fox14, McDermott11, DelNobile13a, DelNobile13b, Cirelli13, NRopsDD, DelNobile14a, Feldstein14, Cherry14} would be useful for e.g.~giving exclusion limits from or comparing sensitivities of different detectors/experiments. }, one minor contribution of the \amidas\ package to direct DM detection experiments would be to be helpful in future detector design and material search. Firstly, by running Monte Carlo simulations with \amidas\ one can estimate the required experimental exposures for reconstructing WIMP properties with acceptable statistical uncertainties and/or systematic biases. Secondly and more importantly, it has been discussed that, for determining the WIMP mass one frequently used heavy target nucleus needs to be combined with a light one \cite{DMDDmchi}; for determining ratios between different WIMP--nucleon cross sections not only the conventionally used targets but also those with non--zero expectation values of both of the proton and neutron group spins are required \cite{DMDDranap}. In addition, different combinations of chosen detector materials with different nuclear masses and/or total nuclear spins and (ratios between) the nucleon group spins not only can affect the statistical uncertainties \cite{DMDDmchi, DMDDranap} but also are available in different range of interest of the reconstructed properties \cite{DMDDranap}. In this paper, we give a detailed user's guide to the \amidas\ package and website. In Sec.~2, we list the \amidas\ functions and different simulation and data analysis modes for these functions. The use of intrinsically saved element data to set information of users' favorite target nuclei will be described in detail. In Sec.~3, we describe the meanings and options of all input parameters/factors for running Monte Carlo simulations. The preparation of uploaded files for defining the one--dimensional WIMP velocity distribution function, the elastic nuclear form factors for SI and SD WIMP--nucleus cross sections as well as artificial/experimental background spectrum will be particularly described. In Sec.~4, the preparation of the real/pseudo data files and the uploading/analyzing procedure on the \amidas\ website will be given. The new developed Bayesian analysis technique \cite{DMDDf1v-Bayesian} will be talked separately in Sec.~5. We conclude in Sec.~6. Some technical detail in the \amidas\ package will be given in Appendix. | In this paper, we give a detailed user's guide to the \amidas\ (A Model--Independent Data Analysis System) package and website, which is developed for online simulations and data analyses for direct Dark Matter detection experiments and phenomenology. \amidas\ has the ability to do full Monte Carlo simulations as well as to analyze real/pseudo data sets either generated by another event generating programs % or recorded in direct DM detection experiments. Recently, the whole \amidas\ package and website system has been upgraded to the second phase: \amidasii, for including the new developed Bayesian analysis technique. Users can run all functions and adopt the default input setup used in our earlier works \cite{DMDDf1v, DMDDf1v-Bayesian, DMDDmchi, DMDDfp2, DMDDranap, DMDDbg-mchi, DMDDbg-f1v, DMDDbg-fp2, DMDDbg-ranap, AMIDAS-CYGNUS2011, AMIDASbg-DSU2011} for their simulations as well as analyzing their own real/pseudo data sets. The use of the \amidas\ website for users' simulations and data analyses has been explained step--by--step with plots in this paper. The preparations of function/data files to upload for simulations and data analyses have also been described. Moreover, for more flexible and user--oriented use, users have the option to set their own target nuclei as well as their favorite/needed (fitting) one--dimensional WIMP velocity distribution function, elastic nuclear form factors for the SI and SD WIMP--nucleus cross sections and different probability distribution functions needed in the Bayesian reconstruction procedure. As examples, the \amidasii\ codes for all user--uploadable functions are given in Secs.~3 and 5 as well as Appendix B and C. In summary, up to now all basic functions of the \amidas\ package and website have been well established. Hopefully this new tool can help our theoretical as well as experimental colleagues to understand properties of halo WIMPs, offer useful information to indirect DM detection as well as collider experiments, and finally discover (the mystery of) Galactic DM particles. \subsubsection* | 14 | 3 | 1403.5611 |
1403 | 1403.2388_arXiv.txt | {To explore the complex halo substructure that has been reported in the direction of the Virgo constellation, radial velocities and metallicities have been measured for 82 RR Lyrae stars (RRLS) that were identified by the QUEST survey. These stars are distributed over 90 sq. deg. of the sky, and lie from 4 to 23 kpc from the Sun. Using an algorithm for finding groups in phase space and modeling the smooth halo component in the region, we identified the 5 most significant RRLS groups, some of which were previously known or suspected. We have examined the SEKBO and the Catalina catalog of RRLS (with available spectroscopic measurements by Prior et al. 2009, and Drake et al. 2013), as well as the bright QUEST RRLS sample (Vivas et al. in prep.), the catalogs of blue horizontal branch (BHB) stars compiled by Sirko et al (2004) and Brown et al (2008, 2010) and the catalog of Red Giant stars from the Spaghetti survey, for stars that may be related to the QUEST RRLS groups. The most significant group of RRLS is the Virgo Stellar Stream (VSS, first reported by Duffau et al 2006) identified here as group A, which is composed of at least 10 RRLS and 3 BHB stars. It has a mean distance of 19.6 kpc and a mean radial velocity $V_{\rm gsr} = 128$ ${\rm km~s}^{-1}$, as estimated from its RRLS members. With the revised velocities reported here, there is no longer an offset in velocity between the RRLS in the VSS and the prominent peak in the velocities of main-sequence turnoff stars reported by Newberg et al (2007) in the same direction and at a similar distance (S297+63-20.5). The location in phase space of two other groups (F and H) suggests a possible connection with the VSS, which cannot be discarded at this point, although the turnoff colors of the VSS and group H, as identified from Newberg et al. 2007, suggest they might be composed of different populations. Two more groups, B and D, are found at mean distances of 19.0 and 5.7 kpc, and mean radial velocities of $V_{\rm gsr} = -94$ and $32$ ${\rm km~s}^{-1}$. The latter is the more numerous in terms of total members, as well as the more extended in RA. A comparison with the latest model of the disruption of the Sagittarius dwarf, indicates that none of the above groups is related to it. Rather than being the result of a single accretion event, the excess of stars observed in Virgo appears to be composed of several halo substructures along the same line of sight. } | } Over the last decade, a number of simulations of galaxy formation that are based on the popular $\Lambda$CDM hierarchical picture have predicted that the halos of disk galaxies should contain numerous substructures that are the debris from accreted dwarf galaxies \citep{bullock05,cooper10}. Over the same period, the evidence that the Milky Way has indeed accreted dwarf galaxies has grown considerably \citep[e.g.][]{bell08,helmi11,xue11}. While there is no doubt that accretion has occurred, it is not yet firm that the number and the properties of the observed substructures are consistent with those simulations. To answer this question, it is not only necessary to find the substructures but also to characterize them in sufficient detail that the masses of the dwarf galaxies and the times of their accretion can be estimated. This paper reports a more thorough description of halo substructure in the direction of the constellation Virgo. After the Sagittarius (Sgr) tidal streams \citep[e.g.][]{majewski03}, the Virgo region contains the most obvious overdensity of stars in the sky explored so far by large scale surveys. Discovered as an overdensity of RR Lyrae stars (RRLS) in the QUEST survey \citep{vivas01,vivas02,vivas06,duffau06} and of main-sequence turn-off stars \citep{newberg02} in the SDSS, it is one of the most noticeable features in the Field of Streams \citep{belokurov06}. Based in main sequence stars from SDSS, \citet{juric08} estimated that the substructure, covers $\sim 1000$ sq deg of the sky. More recently, \citet{bonaca12} suggested the feature may span up to $\sim 3000$ sq deg. \citet{juric08} suggested that the Virgo overdensity (VOD) was produced by the merger of a low surface brightness galaxy with the Milky Way (see also \citet{carlin12}). Spectroscopic observations have revealed, however, a complex system of substructures in the kinematic distribution of stars \citep{duffau06,newberg07,vivas08,starkenburg09,prior09a,brink10,casey12}. Furthermore, it appears that the strengths of these features depend on the tracer used. For example, the strongest peak in the velocity distribution of RRLS and turnoff stars in the region is found near $V_{\rm gsr}=120$ ${\rm km~s}^{-1}$ \citep{duffau06,newberg07,prior09a}, while main sequence and K giants show a peak at $V_{\rm gsr} = -80$ ${\rm km~s}^{-1}$ \citep{brink10,casey12}. In their study of substructures in high resolution simulations of galactic halos from the Aquarius project, \citet{helmi11} found that debris from massive progenitors often look like diffuse substructures on the sky. They point out that because of the large-scale structure present in the Universe when galaxies began to form, the accretion of galaxies was not random in direction, but was instead along preferred directions. This infall pattern causes the streams and substructures of different progenitors to overlap. This scenario, rather than a single massive progenitor, may be the explanation for the substructures in Virgo. To complicate the picture, the region of the VOD is not far from the leading tail of the Sgr dSph galaxy. Although the bulk of the Sgr stars lie much farther away, at 50 kpc, it has been suggested that stars which became unbound in previous passages of the galaxy may lie at much closer distances \citep{martinez07,prior09a,law10}. Although the distance to the Virgo overdensity is relatively small, the spatial distribution of stars alone does not provide a clear picture of the complex substructure. The densest part of the RRLS overdensity lies at 19 kpc \citep{vivas06}, but \citet{juric08} suggested that the VOD covered a range of distances from $\sim 5$ to 20 kpc. Velocities are key to separate the different accretion events and establish the relationship between them, if any. In this investigation, we obtained spectroscopy of RRLS present in the QUEST survey in the Virgo region between 4 and 23 kpc, with the goal of filling the gap between the earlier investigations by \citet{duffau06} and \citet{vivas08} who observed QUEST RRLS in ranges 18-20 kpc and $< 12$ kpc, respectively. The sample of RRLS presented here, which is a combination of new observations and an updating of previous ones, includes 87\% of the QUEST RRLS in an area of almost 90 sq degrees of the sky. The relative errors of the RRLS distances are about 7\% (Vivas and Zinn 2006), which is superior to most other halo tracers (e.g., $\sim 15\%$ for red giants, Starkenburg 2009) and comparable to that obtained for BHB stars, their nearest rival. The pulsations of RRLS make the determination of their systemic radial velocities more difficult than for other stars, but even using small numbers of observations of modest precision, it is possible to obtain precisions of $\sim15-20$ ${\rm km~s}^{-1}$ (e.g. Layden 1994, Vivas et al. 2005, Prior et al. 2009). Since the line-of-sight (los) velocity dispersion of the halo is $\sim 115$ ${\rm km~s}^{-1}$ over the distance range considered here \citet{brown10}, this precision is adequate for identifying the velocity peaks produced by substructure against the background of random halo stars. The main reason to use RRLS as tracers is that even when they do not represent the most numerous population of a system, they provide precise distances and good velocities to help pin-point the location in the sky where to continue searching for evidence of an accretion event. The few RRLS and BHB stars one might find clustered in a particular place in the sky are thus highly significant. The confusion caused by the use of a more numerous tracer with a larger distance error in regions where several stream candidates are suspected can be clarified by the use of RRLS. The trade off is that the number of stars found will be smaller but their parameters will be very reliable. This paper is organized as follows: in Section~\ref{sec-sample} we describe the sample of RRLS observed here and its relationship (in space) with other works carried out in the region. Section~\ref{sec-observations} explains the observations and processing techniques that we have used. We look for coherent groups in the sample of RRLS in Section~\ref{sec-results}, and investigate whether or not these groups are also found in surveys using different tracers. Finally, we discuss our results in Section~\ref{sec-conclusions}. | } We present the first complete spectroscopic study of all RRLS in the QUEST survey in the direction of the VSS. For this work we have gathered 82 targets combining two previous works by the collaboration \citep{duffau06,vivas08}, and reporting on 36 RRLS newly observed stars. The sample covers heliocentric distances between 4-23 kpc, $\sim 90$ square degrees, and RA between $178\degr$-$200\degr$. This homogeneous sample enables us to disentangle in phase space several substructures lying along the same line of sight. This is possible because RRLS are excellent standard candles and allow the determination of precise distances. A summary of all kinematic groups found in the Virgo region is displayed in Table~\ref{comp-tab}. In this table, each check mark ($\surd$) under a study reference represents a detection of the corresponding group. The corresponding references to each study are: N07 \citep{newberg07}, B10 \citep{brink10}, C12 \citep{casey12}, S09 \citep{starkenburg09} and P09 \citep{prior09a}. The first column contains, in addition to the 5 groups of RRLS detected here, two groups predicted to lie in this region by the models of the disruption of Sgr (``Sgr far'' refers to stars located at distances $>40$ kpc, while ``Sgr near'' correspond to stars between $\sim 8$ and 18 kpc, as shown by the debris (asterisks) in Figure~\ref{fig-sgr}). An asterisk in Table~\ref{comp-tab} indicates a weak detection of the corresponding group. \begin{table*} \centering \caption{Kinematic Detections in the Virgo Region} \label{comp-tab} \begin{tabular}{lcccccc} \hline \hline Group & This work & N07 & B10 & C12 & S09 & P09 \\ & RRLS & F-TO & F-TO & K-giants & K-giants & RRLS \\ \hline Group A (VSS) & $\surd$ & $\surd$ & $\surd$ & $\surd$ & & $\surd$ \\ Group B & $\surd$ & * & $\surd$ & & & \\ Group D & $\surd$ & & & & & \\ Group F & $\surd$ & & $\surd$ & $\surd$ & & \\ Group H & $\surd$ & $\surd$ & & & $\surd$ & \\ Sgr far & & & $\surd$ & $\surd$ & & \\ Sgr near & & $\surd$ & * & & $\surd$ & $\surd$ \\ \hline \end{tabular} \end{table*} The VSS, which is now reconciled with feature S297+63-20.5 \citep{newberg07}, is the most significant feature in our sample. It spans at least between 17.9-21.6 kpc along the line of sight and has a mean radial velocity of 128 ${\rm km~s}^{-1}$, as deduced from its RRLS components. A possible relationship of the VSS with two other groups (Groups F and H) cannot be discarded at this point. The three groups form a sequence in phase space with the VSS and group F at the two ends of the sequence with group H intermediate in both velocity and distance (see Figure~\ref{fig-groups_extra} and Figure~\ref{fig-sgr}). Another significant group, composed mostly of type c stars, was found at a similar distance than the VSS, 19 kpc, but with a velocity of $-94$ ${\rm km~s}^{-1}$. Similar velocity signatures have been found in the region by \citep{brink10} and \citet{newberg07}, although the latter found it among the brightest stars in their sample, hence, presumably corresponding to a group at a shorter distance than our detection. Finally, a nearby group at only 6 kpc was also detected with a mean velocity of $V_{\rm gsr}=32$ ${\rm km~s}^{-1}$. The properties of this group are harder to isolate due to its much larger angular extension in the sky and its velocity being close to the expected mean velocity of halo stars. Using the models of \citet{law10}, we find that while the Sgr stream debris are not expected to be dense in this region at distances closer than 40 kpc, some debris is expected to be found in the VSS area. Several of our RRLS lie within the expected distance and velocity of those predictions. However, none of the kinematic groups we found in this work seem to be related with Sgr debris. For the distances involved in this study the dominant feature is the VSS. This was not the case in the works of \citet{brink10} and \citet{casey12}. Their exploration included targets presumably at much larger distances. In those cases, the dominant feature seems to be the Sgr leading stream or another new sub-structure in the region as suggested by \citet{jerjen13}. The numerous kinematic features found in our sample of RRLS qualitatively agree with the expectations of halo substructures in cosmological models of galaxy formation. Our results suggest that several streams, rather than a single accretion event, are responsible for the excess of stars found in the Virgo region. The true shape and extension of all the streams detected in this work will require accurate distances and velocities of targets in a more extended region than the one considered here. The possibility of clarifying which of the different substructures are related to each other and which are not, might require additional information. One possibility is the chemical tagging of their stars, which is one direction we are currently pursuing. | 14 | 3 | 1403.2388 |
1403 | 1403.5561_arXiv.txt | The individual star formation histories of bulges and discs of lenticular (S0) galaxies can provide information on the processes involved in the quenching of their star formation and subsequent transformation from spirals. In order to study this transformation in dense environments, we have decomposed long-slit spectroscopic observations of a sample of 21 S0s from the Virgo Cluster to produce one-dimensional spectra representing purely the bulge and disc light for each galaxy. Analysis of the Lick indices within these spectra reveals that the bulges contain consistently younger and more metal-rich stellar populations than their surrounding discs, implying that the final episode of star formation within S0s occurs in their central regions. Analysis of the $\alpha$-element abundances in these components further presents a picture in which the final episode of star formation in the bulge is fueled using gas that has previously been chemically enriched in the disc, indicating the sequence of events in the transformation of these galaxies. Systems in which star formation in the disk was spread over a longer period contain bulges in which the final episode of star formation occurred more recently, as one might expect for an approximately coeval population in which the transformation from spiral to S0 occurred at different times. With data of this quality and the new analysis method deployed here, we can begin to describe this process in a quantitative manner for the first time. | \label{sec:introduction} Lenticular galaxies (S0s) lie between spirals and ellipticals on the Hubble Sequence, sharing the discy morphology of the spiral galaxies but containing the redder colours and old stellar populations generally seen in ellipticals. As a result, they are often seen as a transitional phase between spirals and ellipticals, and so understanding their formation is key to understanding the evolution of galaxies and the significance of the Hubble Sequence. Plenty of evidence exists to suggest an evolutionary link between spirals and S0s. The morphology--density relation of \citet{Dressler_1980} showed that spirals tend to dominate lower density regions of space, such as the field, while S0s begin to dominate as you move into groups and clusters. This finding suggests that the local environment plays a key role in the transformation of spirals to S0s, where the higher frequency of interactions in groups and clusters contribute significantly to the quenching of star formation in the progenitor spirals. Studies such as \citet{Dressler_1997}, \citet{Fasano_2000} and \citet{Desai_2007} have also found a link with redshift, where the fraction of S0s increased toward lower redshifts while that of spirals decreased, thus revealing that S0s are more common at more recent epochs than in the past. Many processes have been suggested to explain this transformation, most of which focus on an interaction that quenches the star formation in the disc followed by passive evolution as the galaxy fades to an S0. The interaction that triggers the transformation could be with the intra-cluster medium as the galaxy travels through the cluster, where the disc gas is removed by ram pressure stripping \citep{Gunn_1972} or the halo gas is stripped and the star formation quenched by starvation \citep{Larson_1980,Bekki_2002}. Alternatively, interactions with neighbouring galaxies could lead to the gas being stripped by harassment \citep{Moore_1996,Moore_1998,Moore_1999}, or minor mergers could initiate starbursts that use up all of the remaining gas throughout the entire disc \citep{Mihos_1994}. It is still uncertain whether any one of these processes dominate the transformation, or if the process changes with time, environment or luminosity. Evidence that such variations do occur was found by \citet{Barway_2007,Barway_2009}, who found that the transformation process is dependent on the luminosity of the galaxy, where fainter S0s evolved through secular processes while brighter galaxies were created by more turbulent ones. The different processes that have been suggested to explain the transformation of spirals to S0s would affect the bulges and discs in different ways, making their individual star formation histories key to understanding the transformation process. For this reason, many studies have tried looking at variations in the stellar populations of S0s between the bulge and disc for clues to the transformation process that created that galaxy. One way of looking at the stellar populations is to use multi-waveband photometry to measure colour gradients over the galaxy, which provides information on the stellar populations as younger and lower-metallicity stars tend to emit bluer light. \citet{Terndrup_1994}, \citet{Peletier_1996} and \citet{Mollenhoff_2004} all found evidence of negative colour gradients within the bulges of S0s and spirals, suggesting that redder light, and therefore older or more metal-rich stellar populations, are more centrally concentrated within these galaxies. Similarly, \citet{Bothun_1990}, \citet{Peletier_1996} and \citet{Hudson_2010} found that the discs of S0s and spirals are bluer than the bulges, suggesting that disc galaxies either experienced more recent star formation at larger radii \citep{deJong_1996_2}, or have higher metallicities in their nuclear regions \citep{Beckman_1996,Pompei_1997}. Spectroscopic studies of galaxies have also been used to study stellar population gradients across the bulges and discs of S0s. For example, negative metallicity gradients and positive age gradients have been detected in S0s by \citet{Fisher_1996}, \citet{Bell_2000}, \citet{Prochaska_2011} and \citet{Bedregal_2011}, which indicate that the central regions of S0s contain younger and more metal rich stars. Further evidence of recent star formation in bulge regions of S0s has been detected by \citet{Poggianti_2001}, \citet{Ferrarese_2006}, \citet{SilChenko_2006b} and \citet{Kuntschner_2006}, and a study by \citet{Pracy_2013} found evidence of strong positive age gradients within the central $\sim$~1~kpc of the bulges of `k+a' galaxies, which are thought to be a transitional phase between spirals and S0s. Another recent study of `k+a' galaxies by \citet{Rodriguez_2014} also found evidence that the most recent star formation activity in these galaxies was centrally concentrated within the disc, and that the transformation from spirals most likely arose through gentler processes such as ram-pressure stripping or galaxy-galaxy interactions. Such studies of the star formation histories of S0 bulges and discs have revealed age and metallicity gradients across the galaxies, but fail to provide information on whether it represents a gradient within the individual components, or whether it arises simply from the superposition of varying amounts of bulge and disc light, where each component contains stellar populations of distinct different ages and metallicities. To overcome these limitations, we have developed a new method for spectroscopic bulge--disc decomposition \citep{Johnston_2012}, in which a high-quality spectrum of a galaxy is cleanly separated into bulge and disc components wavelength-by-wavelength to create two, one-dimensional spectra representing purely the bulge and disc light. These clean spectra can then be analyzed to determine the ages and metallicities of the bulge and disc with minimal contamination in order to determine the sequence of star-formation events that led to the formation of the S0. In this paper, we set out to analyze the bulge and disc star formation histories spectroscopically for a sample of S0s from the Virgo Cluster, in order to determine the process that triggered their transformation from spirals. The Virgo Cluster was selected as the closest single system with sufficient members to undertake a systematic study of this transformation process. Section~\ref{sec:Observations and Data Reduction} describes the data set and reduction, and Section~\ref{sec:Spectroscopic decomposition} summarizes the method. The results for the stellar populations analysis, the star formation timescales and chemical enrichment are discussed in Section~\ref{sec:Stellar Population Analysis}. The implications of these results for the likely evolutionary tracks followed by S0s are discussed in Section~\ref{sec:Summary}. | \label{sec:Summary} In this paper, by decomposing Virgo Cluster S0 galaxies into clean disc and bulge spectra, we have been able to uncover a number of new facts about these individual components as well as the connections between them. From these data, a coherent quantitative picture of each S0's star-formation history is beginning to emerge, which we summarize in cartoon form in Fig.~\ref{SF_history}. The galaxy starts out as a normal spiral, with an old bulge surrounded by young, star-forming disc. At some traumatic point in the galaxy's life, the gas in the disc is stripped, thus quenching the star formation there, and in the process some of the gas gets dumped in the centre of the galaxy leading to a final burst of star formation in the bulge. The galaxy then fades to the S0 that we see today with a predominantly younger and more metal rich bulge surrounded by an older and more metal poor disc, as so clearly found in Fig.~\ref{Age-Metallicity_Virgo}. Although strong indications of this phenomenon have been found previously through radial variations in age and metallicity in S0 galaxies, this study confirms that the phenomenon can be traced to the superposition of distinct bulge and disc components rather than more general gradients within those components. A subtler probe of the star formation histories of bulges and discs is provided by their $\alpha$-element abundances. As we saw in Fig.~\ref{elemental_abundances2}, there is a significant correlation between Mgb/$\langle$Fe$\rangle$ and age in the bulges of these galaxies, but not in their discs, which can both be understood in the context of Fig.~\ref{SF_history}. The emission from the bulge is dominated by the younger stars from the final burst of star formation, so the value of Mgb/$\langle$Fe$\rangle$ is largely dictated by the gas from which this burst formed, which in this picture originated in the disc and was dumped into the bulge when the galaxy transformed. Thus, it reflects the properties of the gas in the disc at the end of its star-forming life. In general, the longer ago this transformation occurred (and hence the older the age inferred for the bulge), the shorter the star-forming lifetime of the disc because $\tau_{SF (disc)} + t_{\rm bulge}$ in Fig.~\ref{SF_history} reflect the total age of the galaxy. If $\tau_{SF (disc)}$ is relatively short (so $t_{\rm bulge}$ is relatively long), the gas left at the end of the disc's star-forming lifetime will not be so polluted by Fe from type~Ia SNe, so Mgb/$\langle$Fe$\rangle$ will be relatively large, explaining the correlation seen. In the disc, on the other hand, the observed value of Mgb/$\langle$Fe$\rangle$ reflects the more extended and potentially complicated complete star-formation history of this component, as its light will not be dominated by a single star-formation event, and the derived luminosity-weighted age will be similarly complex, so the absence of any correlation in this component is not a surprise. This connection between the polluted gas from the disc and the visible last burst of star formation in the bulge is underlined by Fig,~\ref{elemental_abundances}, which shows the general trend that the Mgb/$\langle$Fe$\rangle$ in the two components are correlated, but that the disc is less Fe-enriched than the bulge. This difference arises because the disc's value for Mgb/$\langle$Fe$\rangle$ reflects its entire star-formation history, some of which will have occurred at early times before the Type~Ia SNe started producing large quantities of Fe, whereas the bulge population is dominated by stars produced from the most polluted disc gas, which will be significantly more Fe enriched. There is also an interesting hint in this figure that the most massive galaxies seem to show the least difference between Mgb/$\langle$Fe$\rangle$ for discs and bulges, which would suggest an earlier transformation leading to less difference in the degree of Fe enrichment, as perhaps a new example of the ``downsizing'' phenomenon. \begin{figure} \includegraphics[width=1\linewidth]{Fig8.eps} \caption{A simplified star formation history for the bulge and disc of an S0 galaxy, showing the relationship between the age of the bulge stellar populations, t$_{bulge}$, and the star formation timescale of the disc, $\tau_{SF (disc)}$. The disc (dot-dash line) experiences continuous star formation, the rate of which declines gradually with time, until the quenching process begins, which finishes soon after when a central star-formation event uses up the remaining disc gas to produce the predominantly-young bulge of the final S0 (solid line). \label{SF_history}} \end{figure} As a final illustration of the physics that underlies Fig.~\ref{SF_history}, Fig.~\ref{BulgeAge_DiscMgFe} shows the clear correlation between Mgb/$\langle$Fe$\rangle$ for the disc component and the age of the bulge. Again, this fits with the finite time available for galaxy evolution, such that if the transformation occurs later then the disc will have had time to become strongly polluted by Fe, reducing Mgb/$\langle$Fe$\rangle$, and the bulge will have undergone its final burst of star formation relatively recently, decreasing its luminosity-weighted measured age. As this discussion indicates, there is now a wealth of information that can be gleaned by decomposing spectra of S0 galaxies into their bulge and disc contributions, in studying the detailed stellar population properties of these individual components. We are at the point of being able not only to put together the general picture of the quenching of disc star formation accompanied by a final episode of bulge star formation shown in Fig.~\ref{SF_history}, but also looking at the variations from galaxy to galaxy to tie down the different histories that different galaxies have witnessed. Clearly, such decompositions would be more robust if carried out in two dimensions using integral field unit (IFU) data rather than long-slit spectroscopy, and the up-coming very large IFU Mapping Nearby Galaxies at APO (MaNGA) survey in SDSS-IV promises the size of sample that will answer the remaining questions about the relative importance of different transformation mechanisms. | 14 | 3 | 1403.5561 |
1403 | 1403.2494_arXiv.txt | {This is the first of a series of investigations into far-IR characteristics of 11 planetary nebulae (PNs) under the {\sl Herschel Space Observatory} Open Time\,1 program, Herschel Planetary Nebula Survey (HerPlaNS).} {Using the HerPlaNS data set, we look into the PN energetics and variations of the physical conditions within the target nebulae. In the present work, we provide an overview of the survey, data acquisition and processing, and resulting data products.} {We perform (1) PACS/SPIRE broadband imaging to determine the spatial distribution of the cold dust component in the target PNs and (2) PACS/SPIRE spectral-energy-distribution (SED) and line spectroscopy to determine the spatial distribution of the gas component in the target PNs.} {For the case of NGC\,6781, the broadband maps confirm the nearly pole-on barrel structure of the amorphous carbon-richdust shell and the surrounding halo having temperatures of 26--40\,K. The PACS/SPIRE multi-position spectra show spatial variations of far-IR lines that reflect the physical stratification of the nebula. We demonstrate that spatially-resolved far-IR line diagnostics yield the ($T_{\rm e}$, $n_{\rm e}$) profiles, from which distributions of ionized, atomic, and molecular gases can be determined. Direct comparison of the dust and gas column mass maps constrained by the HerPlaNS data allows to construct an empirical gas-to-dust mass ratio map, which shows a range of ratios with the median of $195\pm110$. The present analysis yields estimates of the total mass of the shell to be 0.86\,M$_{\odot}$, consisting of 0.54\,M$_{\odot}$ of ionized gas, 0.12\,M$_{\odot}$ of atomic gas, 0.2\,M$_{\odot}$ of molecular gas, and $4\times10^{-3}$\,M$_{\odot}$ of dust grains. These estimates also suggest that the central star of about 1.5\,M$_{\odot}$ initial mass is terminating its PN evolution onto the white dwarf cooling track.} {The HerPlaNS data provide various diagnostics for both the dust and gas components in a spatially-resolved manner. In the forthcoming papers of the HerPlaNS series we will explore the HerPlaNS data set fully for the entire sample of 11 PNs.} | The planetary nebula (PN) phase marks the last throes of stellar evolution for low to intermediate initial mass stars (of about 0.8--8\,M$_{\odot}$, \citealt{kwokbook}). During this phase, the circumstellar envelope of gas and dust, which is created by mass loss in the preceding asymptotic giant branch (AGB) and post-AGB phases, undergoes a dramatic transformation (i.e., ionization, photo-dissociation, and dynamical shaping) caused by the fast wind and the intense radiation from the central star and by the less powerful but often significant interstellar radiation field coming from the surrounding interstellar space. As a consequence, a wide variety of underlying physical conditions are showcased within PNs, from fully ionized hot plasma to dusty cold atomic/molecular clouds, which exist (at least to first order) in a stratified manner around the central star. Therefore, PNs provide excellent astrophysical laboratories to test theories of stellar evolution as well as theories of gas-dust dynamical processes in interacting stellar winds that can also interact with the surrounding interstellar medium (ISM). While PN investigations have been traditionally done through diagnostics of optical emission lines, PNs are bright sources at a wide range of wavelengths from the radio through the UV, and in some cases, even in the X-ray (e.g., \citealt{pottasch84,z89,st01,corradi03,schoen05,sandin08,sahai11,chanplans,gdm13}). Investigations using far-infrared (far-IR) radiation are especially critical to comprehend PNs as complex physical systems in their entirety, because a large fraction of the nebula mass may reside outside the central ionized region (e.g., \citealt{villaver02}). For example, up to about 4\,M$_{\odot}$ of matter has been found in the far-IR halo of \object{NGC\,650} \citep{ueta06,vanhoof13}. However, according to the recent mass budget estimates based on the UV to mid-IR photometric survey of the Magellanic Clouds, the amount of circumstellar dust grains has been severely underestimated: only about 3\% of the ISM dust grains is accounted for in the warm component of the circumstellar envelopes \citep{matsuura09,boyer12}. What this implies is that the most extended cold regions of the circumstellar envelope could contain this ``missing mass'' component, which can only be detected in the wavelength ranges in the far-IR and longer. Recent opportunities provided by the {\sl Spitzer Space Telescope} ({\sl Spitzer\/}; \citealt{werner04}), {\sl AKARI Infrared Astronomy Satellite} ({\sl AKARI\/}; \citealt{murakami07}), and {\sl Herschel Space Observatory} ({\sl Herschel\/}; \citealt{pilbratt10}) have made it possible to probe the very extended, coldest parts of PN haloes at the highest spatial resolutions in the far-IR to date (the beam size of several to a few tens of arcsec; e.g., \citealt{ueta06,su07,vanhoof10,vanhoof13,cox11}). The new far-IR window has not only given access to the bulk of the matter in the farthest reaches of PNs, but also permitted us to probe the interacting boundary regions between the PN haloes and ISM, spawning new insights into the processing of the mass loss ejecta as they merge into the ISM (e.g., \citealt{wareing06,sabin10,zhang12}). Among these recent far-IR opportunities, those provided by {\sl Herschel} are unique: {\sl Herschel} allows simultaneous probing of the multiple phases of the gaseous components in PNs via far-IR ionic, atomic, and molecular line emission. The {\sl Infrared Space Observatory\/} ({\sl ISO\/}; \citealt{iso}) made detections of far-IR lines from about two dozen PNs \citep{liu01} and another two dozen PN progenitors and other evolved stars \citep{fong01,cc01}. However, the {\sl ISO\/} apertures typically covered most of the optically-bright regions of the target objects,\footnote{The aperture size of the {\sl ISO} LWS detector in the spatial dimension is about $106^{\prime\prime}$, while the beam size is about $40^{\prime\prime}$ radius \citep{isohb}.} and therefore, the previous {\sl ISO} spectroscopic analyses were usually performed in a spatially-integrated manner. {\sl Herschel\/}'s spectral mapping capabilities allow us to look for variations of line/continuum strengths as a function of location in the target nebulae, so that the spatially resolved energetics of the circumstellar envelope can be unveiled. Far-IR line maps would help to trace the spatial variations of the electron density, electron temperature, and relative elemental abundance, which may suggest how much of which material was ejected at what time over the course of the progenitor star's mass loss history. Also revealed is how PNs are influenced by the passage of the ionization front. While such line diagnostics have been routinely performed in the optical line diagnostics in the far-IR can offer an alternative perspective, because (1) far-IR line ratios are relatively insensitive to the electron temperature due to smaller excitation energies of fine-structure transitions in the far-IR, and (2) far-IR line and continuum measurements are often extinction-independent, permitting probes into dusty PNs. Hence, PN investigations in the far-IR with {\sl Herschel\/} should have a bearing on abundance determinations and elemental column densities, and therefore can heavily impact analyses in other wavelength regimes. With the foregoing as motivation, we have conducted a comprehensive far-IR imaging and spectroscopic survey of PNs, dubbed the Herschel Planetary Nebula Survey (HerPlaNS), using nearly 200\,hrs of {\sl Herschel\/} time by taking advantage of its mapping capabilities -- broadband and spectral imaging as well as spatio-spectroscopy -- at spatial resolutions made possible by its 3.3\,m effective aperture diameter. Our chief objective is to examine both the dust and gas components of the target PNs simultaneously in the far-IR at high spatial resolutions and investigate the energetics of the entire gas-dust system as a function of location in the nebula. In this first installment of the forthcoming HerPlaNS series of papers, we present an overview of the HerPlaNS survey by focusing on the data products and their potential. Below we will describe the schemes of observations and data reduction (\S\,\ref{obs}), showcase the basic data characteristics using the PN \object{NGC\,6781} as a representative sample (\S\,\ref{6781}), and summarize the potential of the data set (\S\,\ref{sum}) to pave the way for more comprehensive and detailed analyses of the broadband mapping and spectroscopy data that will be presented in the forthcoming papers of the series. | Summary of HerPlaNS Data Products and their Characteristics} \centering \begin{tabular}{lcccccccccc} \hline\hline Observing Mode & Instrument/Band & $\lambda (\Delta\lambda)$ & \multicolumn{2}{c}{Data Characteristics} \\ & & ($\mu$m) & & \\ \hline % Imaging & PACS/Blue & 70 (25) & scan map (5\farcs6 beam at 1$^{\prime\prime}$\,pix$^{-1}$) & $2\farcm5 \times 2\farcm5$ to $7^{\prime} \times 7^{\prime}$ field of view \\ & PACS/Green & 110 (45) & scan map (6\farcs8 beam at 1$^{\prime\prime}$\,pix$^{-1}$) & by 2 orthogonal scans\\ & PACS/Red & 160 (85) & scan map (11\farcs4 beam at 2$^{\prime\prime}$\,pix$^{-1}$) &\\ & SPIRE/PSW & 250 (76) & scan map (18\farcs2 beam at 6$^{\prime\prime}$\,pix$^{-1}$) & $4^{\prime} \times 8^{\prime}$ field of view \\ & SPIRE/PMW & 350 (103) & scan map (24\farcs9 beam at 9$^{\prime\prime}$\,pix$^{-1}$) & by 2 orthogonal scans \\ & SPIRE/PLW & 500 (200) & scan map (36\farcs3 beam at 14$^{\prime\prime}$\,pix$^{-1}$) & \\ Spectroscopy & PACS/B2A & 51--72 & spectral cube ($R\approx4000$ at 9\farcs6\,spaxel$^{-1}$) & $\sim 50^{\prime\prime} \times 50^{\prime\prime}$ field of view \\ & PACS/B2B & 70--105 & spectral cube ($R\approx2000$ at $10^{\prime\prime}$\,spaxel$^{-1}$) & by $5 \times 5$ IFU spaxels\\ & PACS/R1 & 103--145 & spectral cube ($R\approx1500$ at 11\farcs6\,spaxel$^{-1}$)\\ & PACS/R1 & 140--220 & spectral cube ($R\approx1000$ at 13\farcs2\,spaxel$^{-1}$) \\ & SPIRE/SSW & 194--342 & spectral array ($R\approx1000$ at $17^{\prime\prime}$--$21^{\prime\prime}$\,beam$^{-1}$) & $4^{\prime}$ diameter field of view \\ &&&& by a 35-bolometer array \\ & SPIRE/SLW & 316--672 & spectral array ($R\approx500$ at $29^{\prime\prime}$--$42^{\prime\prime}$\,beam$^{-1}$) & $4^{\prime}$ diameter field of view \\ &&&& by a 19-bolometer array \\ \hline \end{tabular} \tablefoot{% See Fig.\,\ref{specposmaps} for relative placements of PACS and SPIRE spectroscopic apertures. The outermost SPIRE bolometers (16 for SSW and 12 for SLW) are located outside of the unvignetted 2\farcm6 field of view.} \end{table*} Here, we briefly summarize the data reduction steps we adopted. Complete accounts of reduction processes will be presented in the forthcoming papers of the series (D.\,Ladjal et al.\ {\sl in prep}; K.\ M.\ Exter et al.\ {\sl in prep}). A summary of the HerPlaNS data products and their characteristics is given in Table\,\ref{datasummary}. \subsubsection{Broadband Imaging} To generate broadband images, we used the Herschel interactive processing environment (HIPE, version 11; \citealt{hipe}) and Scanamorphos data reduction tool (Scanamorphos, version 21; \citealt{scana}). First, the raw scan map data were processed with HIPE from level 0 to level 1. During this stage, basic pipeline reduction steps were applied while the data were corrected for instrumental effects. The level 1 data were then ingested into Scanamorphos, which corrects for brightness drifts and signal jumps caused by electronic instabilities and performs deglitching, flux calibration, and map projection. Scanamorphos was chosen as our map-making engine over other choices -- photoproject (the default HIPE mapper) and MADmap \citep{madmap} -- because it reconstructs surface brightness maps of extended sources with the lowest noise, which is of great importance for our purposes. After processing with HIPE and Scanamorphos, we obtained far-IR surface brightness maps at 5 bands (70, 160, 250, 350, and 500\,$\mu$m) for 11 PNs, each covering at most $7^{\prime} \times 7^{\prime}$ unvignetted field centered at the target source. \subsubsection{PACS Spectroscopy} We used HIPE track 11 with the calibration release version 44 to reduce all of the PACS spectroscopy data of HerPlaNS. Within HIPE, we selected the background normalization PACS spectroscopy pipeline script for long range and SEDs to reduce the range scan data and the same pipeline script for line scans to reduce the line scan data. Our reduction steps follow those described in the PACS Data Reduction Guide: Spectroscopy.\footnote{% \url{http://herschel.esac.esa.int/hcss-doc-9.0/load/pacs_spec/html/pacs_spec.html} (Version 1, Aug.\ 2012)} In the range scan mode, we used the blue bands B2A (51--72\,$\mu$m) and B2B (70--105\,$\mu$m) and each time we also got simultaneous spectra in the R1 band (103--145\,$\mu$m and 140--220\,$\mu$m), achieving the full spectral coverage from 51--220\,$\mu$m. Each observation results in simultaneous spatial coverage of a $\sim50^{\prime\prime} \times 50^{\prime\prime}$ field by a set of $5 \times 5$ spaxels of the IFU (each spaxel covering roughly a $10^{\prime\prime} \times 10^{\prime\prime}$ field). The PACS IFU $5\times5$ data cubes can also be integrated over a specific wavelength range to generate a 2-D line map. This process can be done for any line detected at a reasonable S/N. \subsubsection{SPIRE Spectroscopy} We used the standard HIPE-SPIRE spectroscopy data reduction pipeline for the single-pointing mode (version 11 with SPIRE calibration tree version 11) to reduce all of the SPIRE spectroscopy data of HerPlaNS, but with the following three major modifications; (1) we extracted and reduced signal from each bolometer individually instead of signal from only the central bolometer as nominally done for single-pointing observations; (2) we applied the extended source flux calibration correction to our data; and (3) we used our own dedicated off-target sky observations for the background subtraction (Fig.\,\ref{6781example}). Besides these extra steps, our reduction steps basically copy those described in the SPIRE Data Reduction Guide.\footnote{% \url{http://herschel.esac.esa.int/hcss-doc-9.0/load/spire_drg/html/spire_drg.html} (version 2.1, Document Number: SPIRE-RAL-DOC 003248, 06 July 2012)} The standard apodization function was applied to the data to minimize the ringing in the instrument line shape wings at the expense of spectral resolution. At the end of these processes, each of the on-source (center and off-center) and off-sky pointings would yield 35 short-band spectra\footnote{In the SLW band, 2 bolometers out of the total of 37 are blind.} from individual hexagonal bolometer positions ($33^{\prime\prime}$ spacing between bolometers) for 194--342\,$\mu$m and 19 long-band spectra from individual hexagonal bolometer positions ($51^{\prime\prime}$ spacing between bolometers) for 316--672\,$\mu$m. The bolometer beams for the short and long band arrays overlap spatially at about a dozen positions, from which the full range spectrum (194--672\,$\mu$m) can be constructed. We created an off-sky spectrum by taking a median of spectra taken from the detectors located within the unvignetted field (i.e., all but the outermost bolometers) of each off-sky position and subtracting the off-sky spectrum from each on-source spectrum taken from the unvignetted field of the bolometer array. Data from the vignetted outermost bolometers are not included for the present science analyses because these bolometers are not sufficiently calibrated for their uncertainties and long term stabilities by the instrument team. Because of the large data volume collected and redundant spatial coverage by center and off-center pointings for some of the target sources, it is possible to self-calibrate data from the outermost bolometers. However, this is beyond the scope of the present overview and hence will be discussed in the forthcoming papers of the series. } Using the {\sl Herschel Space Observatory}, we have collected a rich far-IR imaging and spectroscopic data set for a group of 11\,PNs under the framework of the {\sl Herschel} Planetary Nebula Survey (HerPlaNS). In this survey, we used all available observational modes of the PACS and SPIRE instruments aboard {\sl Herschel} to investigate the far-IR characteristics of both the dust and gas components of the circumstellar nebulae of the target sources. We obtained (1) broadband maps of the target sources at five far-IR bands, 70, 160, 250, 350, and 500\,$\mu$m, with rms sensitivities of 0.01--0.1\,mJy\,arcsec$^{-2}$ (0.4--4 MJy\,sr$^{-1}$); (2) $5 \times 5$ IFU spectral cubes of 51--220\,$\mu$m covering a $\sim50^{\prime\prime}\times50^{\prime\prime}$ field at multiple positions in the target sources, with rms sensitivities of 0.1--1\,mJy\,arcsec$^{-2}$ (4--40 MJy\,sr$^{-1}$) per wavelength bin; and (3) sparsely sampled spectral array of 194--672\,$\mu$m covering a $\sim3^{\prime}$ field at multiple positions in the target sources, with rms sensitivities of 0.001--0.1\,mJy\,arcsec$^{-2}$ (0.04--4 MJy\,sr$^{-1}$) per wavelength bin. In this first part of the HerPlaNS series, we described the data acquisition and processing and illustrated the potential of the HerPlaNS data using \object{NGC\,6781}, a dusty molecular-rich bipolar PN oriented at nearly pole-on, as an example. Broadband images unveiled the surface brightness distribution of thermal continuum emission from \object{NGC\,6781}. Spatially resolved was the object's signature ring structure of a $40^{\prime\prime}$ radius with a $20^{\prime\prime}$ width, embedded in an extended halo of about 100$^{\prime\prime}$ radius. This far-IR ring represents a nearly pole-on bipolar/cylindrical barrel structure, containing at least $M_{\rm dust}=4 \times 10^{-3}$\,M$_{\odot}$, at the adopted distance of $950\pm143$\,pc (\citealt{sm06}; all distance-dependent quantities were based on this value and subject to its 15\% uncertainty). Spectral fitting of the broadband images indicated that dust grains are composed mostly of amorphous-carbon based material (i.e., the power-law emissivity index of $\beta \approx 1$) of the temperature between 26 and 40\,K. In the past, far-IR SED fitting with broadband fluxes were performed under the assumption of negligible line contamination. With the {\sl HerPlaNS} data, we verified that the degree of line contamination is approximately 8--20\% and does not significantly affect the fitting results. The {\sl Herschel} spectra obtained at various locations within \object{NGC\,6781} revealed both the physical and chemical nature of the nebula. The spectra revealed a number of ionic and atomic lines such as [\ion{O}{III}]\,52, 88\,$\mu$m, [\ion{N}{III}]\,57\,$\mu$m, [\ion{N}{II}]\,122, 205\,$\mu$m, [\ion{C}{II}]\,158\,$\mu$m, and [\ion{O}{I}]\,63, 146\,$\mu$m, as well as various molecular lines, in particular, high-J CO rotational transitions, OH, and OH$^{+}$ emission lines: see \citet{aleman13} and \citet{etxaluze13} for more details on the discovery of OH$^{+}$ emission in PNs. Thermal dust continuum emission was also detected in most bands in these deep exposure spectra. Moreover, spectra taken at multiple spatial locations elucidated the spatial variations of the line and continuum emission, which reflect changes of the physical conditions within the nebula projected onto the plane of sky. On average, the relative distributions of emission lines of various nature suggested that the barrel cavity is uniformly highly ionized, with a region of lower ionization delineating the inner surface of the barrel wall, and that the least ionic and atomic gas, molecular, and dust species are concentrated in the cylindrical barrel structure. The CO rotation diagram diagnostics yielded $T_{\rm ex} \approx$\,60--70\,K and $N_{\rm CO}\approx 10^{15}$\,cm$^{-2}$. Compared with the previous CO measurements and diagnostics by \citet{bachiller93}, the present observations and analysis with higher-J transitions sampled much warmer CO gas component in the cylindrical barrel structure, probably located closer to the equatorial region along the line of sight. However, the amount of this warm component was determined to be an order of magnitude smaller than the cold component. Based on the PACS IFU spectral cube data, we derived line maps in the detected ionic and atomic fine-structure lines. Then, diagnostics of the electron temperature and density using line ratios such as [\ion{O}{III}] 52/88\,$\mu$m and [\ion{N}{II}] 122/205\,$\mu$m resulted in ($T_{\rm e}$, $n_{\rm e}$) and ionic/elemental/relative abundance profiles for the first time in the far-IR for any PN. The derived $T_{\rm e}$ profile substantiated the typical assumption of uniform $T_{\rm e} = 10^4$\,K in the main ionized region, while showing an interesting increase in the barrel wall up to 11,000\,K, followed by a sudden tapering off toward the halo region. The $n_{\rm e}$ profile of high-excitation species is nearly flat at $\sim400$\,cm$^{-3}$ across the inner cavity of the nebula, whereas the $n_{\rm e}$ profile of low-excitation species exhibits a radially increasing tendency from 80\,cm$^{-3}$ to $>600$\,cm$^{-3}$ with a somewhat complex variation around the barrel wall. In fact, this $n_{\rm e}$[\ion{N}{II}] profile is reflected in the physical stratification of the nebula revealed by the ionic/elemental abundance analysis. We found (1) a very highly ionized, centrally restricted \ion{H}{II}R within $\sim20^{\prime\prime}$, (2) a highly ionized \ion{H}{II}R for 20--30$^{\prime\prime}$ of the center marked by a high relative abundance of N$^{2+}$ and O$^{2+}$, (3) a moderately ionized \ion{H}{II}R within the inner surface of the barrel wall (for 30--50$^{\prime\prime}$) marked by a high relative abundance of N$^{+}$ and C$^{+}$, and (4) a least ionized \ion{H}{II}R transitioning into a PDR on the barrel wall (beyond $\sim50^{\prime\prime}$) marked by the presence of molecular and dust species. The detected stratification is consistent with the previous inferences made from the past optical imaging observations in various emission lines of varying levels of excitation. The derived relative elemental abundance profiles showed uniformly low N and C abundances, confirming the low initial mass ($<2$\,M$_{\odot}$) and marginally carbon-rich nature of the central star. However, the profiles did not appear to reveal variations reflecting the evolutionary change of the central star, such as a radially increasing carbon abundance. Nevertheless, the range of relative elemental abundances measured for spatially resolved observations of \object{NGC\,6781} overlaps with the range of abundances obtained from spatially-unresolved measurements of ORLs and CELs. This may indicate that the issue of dichotomy between abundance measurements made from ORLs and CELs is simply due to the spatial resolution effects. Therefore, it is interesting to revisit spatially-resolved diagnostics using optical line maps as performed here using far-IR line maps. Direct comparison between the dust column mass map derived from the HerPlaNS broadband thermal dust data and the gas column mass map derived from the HerPlaNS fine-structure line mapping data augmented with literature data in other wavelengths yielded an empirical gas-to-dust mass ratio distribution map for \object{NGC\,6781}. The resulting empirical gas-to-dust mass ratio map showed a range of ratios within the cylindrical barrel structure, in general radially decreasing roughly from 550 to 100. The average gas-to-dust mass ratio in the dust barrel was determined to be $195\pm110$, and hence, is generally consistent with the typical spatially-unresolved ratio of 100--400 widely used in the literature for the case of PNs and AGB stars. The HerPlaNS data would therefore allow further investigations of the distribution of gas-to-dust mass ratios across the target PNs in a spatially resolved manner. The derivation of column mass distribution maps for various components of \object{NGC\,6781}, based on the empirically-established distribution of the ionized gas component, yielded an estimate for the total mass of the shell of 0.86\,M$_{\odot}$, consisting of 0.54\,M$_{\odot}$ of ionized gas, 0.12\,M$_{\odot}$ of atomic gas, 0.2\,M$_{\odot}$ of molecular gas, and $4\times10^{-3}$\,M$_{\odot}$ of dust grains. Provided that the present core mass is 0.6\,M$_{\odot}$, we concluded that the progenitor star had an initial mass of 1.5\,M$_{\odot}$. Then, theoretical evolutionary tracks of this 1.5\,M$_{\odot}$ star would suggest that the star is nearing to the end of its PN evolution, transitioning onto the white dwarf cooling track. In the forthcoming papers of the HerPlaNS series, we will focus on separate analyses of the broadband maps (Ladjal et al.\ {\sl in prep}) and spectra (Exter et al.\ {\sl in prep}) using the wealth of the entire HerPlaNS data set to present more in-depth results of the analyses outlined above. In addition, we will report more on the energetics of the entire gas-dust system as a function of location in the nebulae, emphasizing the results' statistical implications. | 14 | 3 | 1403.2494 |
1403 | 1403.0491_arXiv.txt | {} { Recent magnetic field surveys in O- and B-type stars revealed that about 10\% of the core-hydrogen-burning massive stars host large-scale magnetic fields. The physical origin of these fields is highly debated. To identify and model the physical processes responsible for the generation of magnetic fields in massive stars, it is important to establish whether magnetic massive stars are found in very young star-forming regions or whether they are formed in close interacting binary systems. } { In the framework of our ESO Large Program, we carried out low-resolution spectropolarimetric observations with FORS\,2 in 2013 April of the three most massive central stars in the Trifid nebula, HD\,164492A, HD\,164492C, and HD\,164492D. These observations indicated a strong longitudinal magnetic field of about 500--600\,G in the poorly studied component HD\,164492C. To confirm this detection, we used HARPS in spectropolarimetric mode on two consecutive nights in 2013 June. } { Our HARPS observations confirmed the longitudinal magnetic field in HD\,164492C. Furthermore, the HARPS observations revealed that HD\,164492C cannot be considered as a single star as it possesses one or two companions. The spectral appearance indicates that the primary is most likely of spectral type B1--B1.5\,V. Since in both observing nights most spectral lines appear blended, it is currently unclear which components are magnetic. Long-term monitoring using high-resolution spectropolarimetry is necessary to separate the contribution of each component to the magnetic signal. Given the location of the system HD\,164492C in one of the youngest star formation regions, this system can be considered as a Rosetta Stone for our understanding of the origin of magnetic fields in massive stars. } {} | Magnetic fields have fundamental effects not only on the evolution of massive stars, on their rotation, and on the structure, dynamics, and heating of their radiatively-driven winds, but also on their final display as supernova or gamma-ray burst. About a few dozen massive magnetic stars are currently known, just enough to establish the fraction of magnetic, core-hydrogen burning stars to be of the order of 8\% (Grunhut et al.\ \cite{grunhut2012}), which appears to be similar to that of intermediate-mass stars. While it is established that the magnetic fields in massive stars are not dynamo-supported, but stable with decay times exceeding the stellar lifetime, their origin is highly debated. The two main competing ideas are that the fields are either ``fossil'' remnants of the Galactic ISM field that are amplified during the collapse of a magnetised gas cloud (e.g.\ Price \& Bate\ \cite{price2007}), or that they are formed in a dramatic close-binary interaction, i.e., in a merger of two stars or a dynamical mass transfer event (e.g.\ Ferrario et al.\ \cite{fer2009}). The intermediate-mass stars show a magnetic fraction during their pre-main sequence evolution similar to Herbig stars (e.g., Hubrig et al.\ \cite{Hubrig2009}; Hubrig et al.\ \cite{Hubrig2013} and references therein; Alecian et al.\ \cite{alecian13}) as their main-sequence descendants, which may speak for fossil fields. On the other hand, there are almost no close binaries amongst the magnetic intermediate-mass main-sequence Ap-type stars (e.g.\ Carrier et al.\ \cite{car2002}). Since this is expected if they were merger products, this argues for the binary hypothesis of the field origin. We present our spectropolarimetric observations of three massive stars in the Trifid Nebula in the framework of our ``B fields in OB stars'' (BOB) collaboration. This nebula is a very young ($\lesssim 10^6$\,yrs) and active site of star formation containing a rich population of young stellar objects (YSOs) and protostars (e.g.\ Cernicharo et al.\ \cite{cernicharo98}). The spectacular, well-known optical \ion{H}{ii} region provides an ideal place for investigating the onset of star birth and triggered star formation. This large nebula is ionised by the O7.5~Vz star HD\,164492A (Sota et al.\ \cite{Sota2013}), which is the central object of the multiple system ADS~10991, containing at least seven components (A to G; Kohoutek et al.\ \cite{koho99}). Due to the faintness of the components HD\,164492B and HD\,164492E-G (all have visual magnitudes fainter than 10.6), we only searched for a magnetic field for the three most massive components, HD\,164492A, C, and D. | \label{sect:disc} Using FORS\,2 and HARPS in the framework of our ESO Large Program 191.D-0255, we detected a magnetic field in the poorly studied system HD\,164492C. Although HD\,164492C appears to be a multiple system, the multiplicity configuration is currently unclear and cannot be better elucidated without additional observations. X-ray emission from HD\,164492C is firmly detected using {\em Chandra} observations, but is blended with a nearby unidentified X-ray source (component C2; Rho et al.\ \cite{rho04}). The total X-ray luminosity of these two marginally spatially resolved sources is $2\times 10^{32}$\,erg\,s$^{-1}$, with both components having similar X-ray brightness. The component C2 shows X-ray variability and is harder in X-rays than HD\,164492C. To identify and model the physical processes that are responsible for the generation of magnetic fields in massive stars, it is important to understand the formation mechanism of magnetic massive stars. Although the Trifid Nebula has often been studied, its distance is not accurately known. Rho et al.\ (\cite{rho08}) reviewed literature values between 1.68 and 2.84\,kpc and adopted a distance of about 1.7\,kpc. Cambr{\'e}sy et al.\ (\cite{cambresy11}) found $2.7\pm0.5$\,kpc in their analysis of new near- and mid-infrared data. Torii et al.\ (\cite{torii11}) used both 1.7 and 2.7\,kpc in their discussion. Even shorter distances of 816\,pc and 1093\,pc were estimated by Kharchenko et al.\ (\cite{kharchenko05,kharchenko13}) by combining proper-motion data with optical and near-infrared photometry in their cluster analysis, respectively. The spatial distribution of the components of the multiple system ADS\,10991 and the photometric study by Kohoutek et al.\ (\cite{koho99}), which revealed almost the same E(B-V) values for components A--C, suggest that these components build a physical system in the nucleus of the Trifid nebula. The age of the Trifid Nebula is only a few 0.1\,Myr according to Cernicharo et al.\ (\cite{cernicharo98}), who considered the spatial extent of the \ion{H}{ii} region. The age of the cluster M20 and the time interval of the star formation in this cluster, of which the system HD\,164492 is a member, can probably be larger, of the order of 1\,Myr. Torii et al.\ (\cite{torii11}) argued that the formation of first-generation stars in the Trifid nebula, including the main ionising O7.5 star HD\,164492A, was triggered by the collision of two molecular clouds on a short time-scale of $\sim$1\,Myr. New insights into the understanding of massive star formation in the Trifid Nebula can be expected from a recently established large project based on infrared and X-ray observations of 20 massive star-forming regions, among them the Trifid Nebula (Feigelson et al.\ \cite{feig2013}). The presented first detection of a magnetic massive multiple system in one of the youngest star-forming regions implies that this system may play a pivotal role in our understanding of the origin of magnetic fields in massive stars. Future spectropolarimetric monitoring of this system is urgently needed to better characterise the components, their orbital parameters, and the magnetic field topology. | 14 | 3 | 1403.0491 |
1403 | 1403.6754_arXiv.txt | It is shown here that a subset of the implicit analytical shock solutions discovered by Becker and by Johnson can be inverted, yielding several exact closed-form solutions of the one-dimensional compressible Navier-Stokes equations for an ideal gas. For a constant dynamic viscosity and thermal conductivity, and at particular values of the shock Mach number, the velocity can be expressed in terms of a polynomial root. For a constant kinematic viscosity, independent of Mach number, the velocity can be expressed in terms of a hyperbolic tangent function. The remaining fluid variables are related to the velocity through simple algebraic expressions. The solutions derived here make excellent verification tests for numerical algorithms, since no source terms in the evolution equations are approximated, and the closed-form expressions are straightforward to implement. The solutions are also of some academic interest as they may provide insight into the non-linear character of the Navier-Stokes equations and may stimulate further analytical developments. | One of the few known non-linear analytical solutions to the equations of fluid dynamics was discovered by \cite{Becker22} and subsequently analyzed by \cite{Thomas44}, \cite{Morduchow49}, \cite{Hayes60} and \cite{Iannelli13}. It captures the physical profile of shock fronts in ideal gases, and although it requires some restrictive assumptions (a steady state, one planar dimension, constant dynamic viscosity, an ideal gas equation of state and a constant Prandtl number $\Pran$ of $3/4$), the solution is exact in the sense that no source terms in the (one-dimensional) evolution equations are neglected or approximated. Analogous solutions were discovered by \cite{Johnson13} in the limit of both large and small $\Pran$. These solutions provide a useful framework for verifying numerical algorithms used to solve the Navier-Stokes equations. A drawback, however, from the perspective of both physical intuition and numerical implementation, is that the solutions are implicit, i.e., they are solutions for $x(v)$ rather than closed-form expressions for $v(x)$ ($x$ here is the spatial dimension in which the shock propagates and $v$ is the velocity magnitude). It is shown here that some of these implicit solutions can be inverted for particular values of the shock Mach number, yielding closed-form expressions for the fluid velocity as a function of position. In particular, for rational values of the shock compression ratio, Becker's implicit expression is a polynomial in $v(x)$. Expressions for the polynomial root relevant to a shock are provided up to a compression ratio of four. Polynomial solutions also exist in both the large- and small-$\Pran$ limits under the assumption of either a constant dynamic viscosity or constant thermal conductivity, and expressions are provided for these as well. Under the assumption of a constant kinematic (rather than dynamic) viscosity, the solution for $v(x)$ takes the particularly simple form of a hyperbolic tangent function; this solution is valid at any Mach number and for both $\Pran \rightarrow \infty$ and $\Pran \rightarrow 3/4$. An overview of the equations to be solved is given in \S\ref{sec:equations}, the solutions are given in \S\ref{sec:solutions}, and a summary is given in \S\ref{sec:summary}. | \label{sec:summary} Several closed-form analytical solutions to the one-dimensional compressible Navier-Stokes equations have been derived in the limit of a steady state and an ideal gas equation of state. Solutions with a constant dynamic viscosity and thermal conductivity can be obtained by solving a polynomial equation. Polynomial solutions valid for large $\Pran$ and $\Pran = 3/4$ are listed in table~\ref{tab:pran34} and shown in figures~\ref{fig:R43_R32}--\ref{fig:R4}. Polynomial solutions valid for small $\Pran$ are listed in table~\ref{tab:pranzero} and shown in figures~\ref{fig:n43_n32}--\ref{fig:nm3_nm2}. Tables~\ref{tab:pran34} and \ref{tab:pranzero} also give expressions for $M_0(\gamma)$ for which these solutions are valid, and the corresponding curves in $M_0$--$\gamma$ space are shown in figure~\ref{fig:M0_gamma}. A solution can also be obtained under the assumption of a constant kinematic viscosity, valid for either large $\Pran$ or a constant $\Pran = 3/4$ and at any Mach number; this solution is described in \S\ref{sec:ckv} and shown in figure~\ref{fig:nu_const}. The derived solutions are non-linear and exact in the sense that no source terms in the evolution equations are neglected or approximated. As such, they make excellent verification tests for numerical algorithms. The most physically relevant solutions are those with $\Pran = 3/4$, as this is close to the $\Pran$ of many gases. The small-$\Pran$ solutions are somewhat relevant to gas mixtures and plasmas, whereas the large-$\Pran$ solutions are primarily of academic interest and are only included for completeness \citep{Johnson13}. The derived solution set is not exhaustive: additional polynomial solutions exist under the assumption of a constant thermal diffusivity $\chi \equiv \kappa/\rho$, and a solution in terms of Lambert functions can be derived for $\mu \propto T^{1/2}$, $\Pran \rightarrow \infty$ and $M_0 \rightarrow \infty$. As none of these solutions are more physically relevant than the ones discussed above, their detailed derivation has not been included. Perhaps the primary benefit of the derived solutions is their addition to the limited number of known exact solutions to the Navier-Stokes equations. Further study of the solutions may provide insight into the non-linear character of these equations, and the methods employed may stimulate additional analytical developments. \begin{figure} \centering \begin{tabular}{cc} \includegraphics[scale=0.4]{figure9a.eps} & \hspace{-0.0in} \includegraphics[scale=0.4]{figure9b.eps} \end{tabular} \caption{Curves in $M_0$--$\gamma$ space for which the derived closed-form solutions are valid, for $R = 4/3$, $3/2$, $2$, $3$ and $4$ (\emph{left, bottom to top}), and for $n = -3$, $-2$, $\infty$, $4/3$, $3/2$, $2$, $3$ and $4$ (\emph{right, top to bottom}). On the right, a dashed line indicates a discontinuous solution, a dotted line indicates a solution with a weak discontinuity, and a solid line indicates a continuous solution.} \label{fig:M0_gamma} \end{figure} I thank the referees for their comments. This work was performed under the auspices of Lawrence Livermore National Security, LLC, (LLNS), under Contract No.$\;$DE-AC52-07NA27344. \appendix | 14 | 3 | 1403.6754 |
1403 | 1403.1032_arXiv.txt | One important strength of the microlensing method in detecting extrasolar planets is its high sensitivity to low-mass planets. However, it is often believed that microlensing detections of Earth-mass planets from ground-based observation would be difficult due to the limit set by finite-source effects. This view comes from the previous estimation of the planet detection probability based on the fractional deviation of planetary signals, but proper probability estimation requires to additionally consider the source brightness, which is directly related to the photometric precision. In this paper, we reevaluate the feasibility of low-mass planet detections considering photometric precision for different populations of source stars. From this, it is found that contribution of the improved photometric precision to the planetary signal of a giant-source event is big enough to compensate the decrease of the magnification excess caused by finite-source effects. As a result, we find that giant-source events are suitable targets for Earth-mass planet detections with significantly higher detection probability than events involved with source stars of smaller radii and predict that Earth-mass planets would be detected by prospective high-cadence surveys. | Searches for extrasolar planets by using the microlensing method is being conducted toward the Galactic bulge field (OGLE: \citet{udalski03}, MOA: \citet{bond01}, \citet{sumi03}, Wise: \citet{shvartzvald12}). Due to the high sensitivity to planets that are difficult to be detected by other methods, the method is important for the comprehensive understanding of the formation and evolution of planets in various types of stars \citep{mao91,gould92,gaudi12}. One of the important strength of the microlensing method is that it is sensitive to low-mass planets. This is because the amplitude of a microlensing planetary signal does not depend on the planet mass for a point source, although the duration of the signal becomes shorter with the decrease of the planet mass. In practice, the low mass limit of a microlensing planet is set by finite-source effects which wash out planetary signals. For giant source stars, the size of the caustic induced by an Earth-mass planet is equivalent to the angular size of the source star and thus the planetary signal is significantly weakened due to severe finite-source effects. For events associated with a main-sequence (MS) star, on the other hand, the attenuation of the planetary signal is mild but poor photometric precision caused by the source faintness limits secure detections from ground-based observation. As a result, it is often believed that detecting Earth-mass planets from ground-based observation would be difficult. The difficulty of detecting Earth-mass planets from ground-based microlensing observation was first pointed out by \citet{bennett96}. The basis of their result lies on their estimation of the planet detection probability based on the simulation of planetary lensing events involved with various types of source stars. In their simulation, they computed the fractional deviation of the lensing magnification, $(A-A_{0})/A_{0}$, and estimated the detection probability by imposing a threshold deviation. Here $A$ and $A_{0}$ represent the lensing magnifications with and without the presence of the planet, respectively. With this criterion, the planet detectability is mostly decided by the severeness of finite-source effects and thus they reached a conclusion that the detection probability of giant-source events would be significantly lower than the probability of events associated with faint source stars with smaller radii. Based on this result, \citet{bennett02} proposed a space-based microlensing experiment in search for Earth-mass planets by resolving faint MS stars. However, proper estimation of the planet detection probability requires to additionally consider the source brightness. This is because the strength of a planetary signal $\Delta\chi^2=\sum_{i}{(F_{i}-F_{0,i})}^2/{\sigma_{i}}^2$ depends not only on the amplitude of the planetary deviation, $F-F_{0}$, but also on the photometric uncertainty, $\sigma$, which is directly related to the source brightness. Here $F$ and $F_{0}$ represent the observed source fluxes with and without the planet, respectively. In addition to the direct decrease of photon noise with the increased photon count, photometric precision further depends on the source brightness because bright source events are likely to be less affected by blending. In this paper, we reevaluate the feasibility of ground-based detections of Earth-mass planets by additionally considering the dependence of photometric precision and blending on the source type. In Section 2, we describe the simulation of planetary lensing events conducted for the estimation of the probability. In Section 3, we present results from the analysis. We summarize the result and conclude in Section 4. | In order to check the previous result indicating the difficulty in detecting Earth-mass planets from ground-based observation, we reevaluated the detection feasibility by additional considering the dependence of photometric precision on source populations. From the analysis based on realistic simulation of lensing events, we found that giant-source events were suitable targets for Earth-mass planet detections with substantially higher detection probability than events involved with source stars of smaller radii. We found the reason for the opposite result to the previous one is that the contribution of the improved photometric precision to the planetary signal of a giant-source event is big enough to compensate the decrease of the magnification excess caused by finite-source effects. Although no Earth-mass planet has been detected yet, the competing effects of extended source and source brightness are well illustrated by microlensing planets detected through the channel of high-magnification events. A good example is the planetary event MOA-2007-BLG-400 \citep{dong09}. Although the mass of the planet $(\sim 0.5 - 1.3$ $M_{\rm Jup})$ discovered in the event is much heavier than the Earth, the event is similar to a giant-source Earth-mass planetary event in the sense that the caustic was substantially smaller than the source and the planetary deviation occurred when the source was bright. Despite that the planetary deviation was greatly attenuated by severe finite-source effects, the signal was detected with the large significance of $\Delta\chi^2 = 1070$ and the planetary deviation was unambiguously ascertained. Considering that the detection was possible mostly thanks to the high photometric precision of the bright source, giant-source events would be suitable targets for the detection of low-mass planets. Therefore, we predict that Earth-mass planets would be detected from future ground-based high-cadence lensing surveys. | 14 | 3 | 1403.1032 |
1403 | 1403.4617_arXiv.txt | We have searched for infrared dust emission from subsets of compact, Galactic neutral hydrogen clouds, with the purpose of looking for dust in high-velocity clouds, identifying low-velocity halo clouds, and investigating the cloud populations defined in the GALFA-HI Compact Cloud Catalog. We do not detect dust emission from high-velocity clouds. The lack of dust emission from a group of low-velocity clouds supports the claim that they are low-velocity halo clouds. We detect dust in the remaining low-velocity clouds, indicating a Galactic origin, with a significantly greater dust-to-gas ratio for clouds with linewidths near 15~\kms. We propose that this is due to dust associated with ionized gas. | The Galactic dust-to-gas emission ratio (DGR) has been observed to be uniform for moderate \HI column densities ($10^{19}-10^{21}$ cm$^{-2}$) at low-velocities ($|$\vlsr$| < 90$\kms) and 100pc scales \citep{boulanger88,jones95}, while the DGR of high-velocity gas is significantly reduced, if not zero \citep{wakker86,miville05,peek09,planck11}. In this paper, we measure the DGR using the ratio of Far-IR flux to \HI column density as traced by 21cm luminosity, assuming optically thin gas. The uniformity of the Galactic DGR is evidence for a well-mixed ISM \citep{boulanger85}. At high \HI column densities, the gas becomes dense enough to self-absorb and to form molecular hydrogen . This reduces the observed 21cm emission which increases the measured DGR for dense regions \citep{blitz90,reach94,reach98,meyerdierks96,boulanger98,douglas07}. Because the Far-IR is tracing the total hydrogen column and not just the \HI, the DGR can be used to detect the presence of 'dark gas' which we cannot detect in 21cm emission \citep{planck11_19}. While high column density gas has an elevated DGR, high-velocity clouds (HVCs) with $|$\vlsr$| > 90$\kms ~have a reduced, if not zero, DGR \citep{wakker86,miville05,peek09}. There are two main explanations for the low DGR, either HVCs have significantly less dust than Galactic gas, or the dust they do have is not radiating sufficiently to be detectable. HVCs are known to have low metallicities, which correlates with a low dust content \citep{wakker97}. HVCs are also known to be far from the Galactic disc, so the interstellar radiation field (ISRF) will not heat the dust as it would the ISM. \cite{peek09} used the reduced DGR of high-velocity clouds (HVCs) to search for their low-velocity analogs - low-velocity halo clouds (LHVCs). LVHCs are an important component of the baryon cycle. If gas is cooling (or staying neutral) in the halo before accreting onto the Galactic disc, simulations show that a significant quantity of gas could be at low radial velocities \citep{sommerLarsen06, peek08}. Other simulations have shown that infalling HVCs could recool after disrupting in the halo \citep{heitsch09}. The recooled clouds would be observable at low radial velocities. Instead of accretion, neutral clouds could condense out of the hot halo \citep{maller04}, but some simulations predict that neutral hydrogen clouds cannot form from linear perturbations \citep{joung12}. Lastly, due to geometric effects, HVCs can have low radial velocities while still having high infall velocities. Detected LVHCs could provide important constraints for models of halo processes (see \cite{putman12} for a review of gas in the Galactic halo). \cite{peek09} identified three candidate LVHCs; two with DGRs agreeing with zero, and one with a very low DGR. Besides searching for LVHCs, the DGR of neutral hydrogen clouds can be used to probe the dust properties of gas that may be part of the Galactic fountain \citep{shapiro76}. Low and intermediate velocity clouds may be in the process of being ejected from the disc or accreting onto it. Recent models by \cite{marasco12} show that supernova-driven accretion could be a major source of fuel for star formation. In this model, ejecta from supernovae provide the nonlinear perturbations necessary for cloud condensation. Observations of ionized metals in the lower halo \citep{shull09,collins09,lehner12} agree with the supernova-driven model \citep{fraternali13} indicating that corresponding neutral clouds could be dust rich. The GALFA-\HI Compact Cloud Catalog \citep{saul12}, hereafter GC3, is an ideal dataset to probe lower column densities, smaller size scales, and search for dust in HVCs or a lack thereof in LVCs. By definition, the clouds in the GC3 are distinct from extended Galactic \HI emission and smaller than 20\arcmin ~in angular size. The GC3 is separated into populations of clouds based on position, velocity, and linewidth distributions that Far-IR detections could support. In this work we make two sets of measurements: the 100\micron/21cm and the 60\micron/21cm ratio in compact clouds. From previous observations, we expect that the 60\micron/21cm ratio will be a factor of four lower than the 100\micron/21cm ratio so we will use the 100\micron/21cm ratio to look for the presence of dust. Where we observe 60\micron ~emission we can explore the size or temperature of the dust grains as smaller grains will radiate more efficiently at shorter wavelengths. Previous observations of the Galactic ISM have measured 100\micron/21cm DGRs of 0.5-2 \dgrunits~and 60\micron/21cm DGRs of 0.1-0.25 \dgrunits~\citep{boulanger88,jones95,reach98,planck11}. We have searched for IR emission from a catalog of small neutral hydrogen clouds to look for answers to several questions: Can we detect dust in HVCs? Is the Galactic DGR constant at smaller scales and lower column densities than have been measured before? Can we detect a population of LVHCs? Does the IR emission agree with the population definitions of \cite{saul12}? This paper is organized as follows: we describe the data we use in \S \ref{sec:data}, in \S \ref{sec:methods} we detail the preprocessing of the data, the DGR detection method, and uncertainty calculations. We report our results in \S \ref{sec:results}, and provide interpretation in \S \ref{sec:discussion}. | \label{sec:conclusion} We have searched for IR dust emission from the compact neutral hydrogen clouds of the GALFA-HI Compact Cloud Catalog using a median image stacking technique and bootstrap error statistics. We use the IRIS reprocessing of the IRAS 100\micron ~and 60\micron ~bands. We do not detect dust in the HVC population and while the uncertainty in our measurements are larger than the previous detections and limits, this is the first constraint for compact clouds as a population. We also do not detect dust in the warm, low-positive velocity clouds in the third Galactic quadrant (+WQ3 clouds), which is additional evidence that these clouds are different from the other LVCs and may be a complex of low-velocity halo clouds. We detect dust in the remaining LVCs, with an elevated DGR for clouds with linewidths of 14-19~\kms. These warm LVCs have a significantly higher DGR than the cold LVCs that may be due to dust associated with the ionized gas component that the cold LVCs may lack. This result is especially striking because we expected to observe elevated dust emission from cold cores in the cold LVCs. The different DGRs of the cold and warm LVCs, along with the difference in the velocity distributions, suggest that these two populations are distinct. We propose that the warm LVCs are associated with the warm, infalling IVC populations, while the cold LVCs are associated with the disc of the Galaxy. Finally, we note that the results of this analysis strongly support the population definitions from \cite{saul12}. | 14 | 3 | 1403.4617 |
1403 | 1403.7223_arXiv.txt | Light echoes, light from a variable source scattered off dust, have been observed for over a century. The recent discovery of light echoes around centuries-old supernovae in the Milky Way and the Large Magellanic Cloud have allowed the spectroscopic characterization of these events, even without contemporaneous photometry and spectroscopy using modern instrumentation. Here we review the recent scientific advances using light echoes of ancient and historic transients, and focus on our latest work on SN~1987A's and Eta~Carinae's light echoes. | Light echoes (LEs) arise when light from a transient or variable source is scattered off circumstellar or interstellar dust, reaching the observer after a time delay resulting from the longer path length \citep[e.g.,][]{Couderc39,Chevalier86,Schaefer87a,Xu94,Sugerman03,Patat05}. Over a century ago, in 1901, the first scattered LEs were discovered around Nova Persei \citep{Ritchey01b,Ritchey01a,Ritchey02}. They were recognized as such shortly thereafter by \citet{Kapteyn02} and \citet{Perrine03}. Since then, LEs have been observed around a wide variety of objects: the Galactic Nova Sagittarii 1936 \citep{Swope40}, the eruptive variable V838 Monocerotis \citep{Bond03}, the Cepheid RS Puppis \citep{Westerlund61,Havlen72}, the T Tauri star S~CrA \citep{Ortiz10}, and the Herbig Ae/Be star R~CrA \citep{Ortiz10}. Echoes have also been observed from extragalactic SNe, with SN~1987A being the most famous case \citep{Crotts88,Suntzeff88}, but also including SNe~1980K \citep{Sugerman12}, 1991T \citep{Schmidt94,Sparks99}, 1993J \citep{Sugerman02,Liu03}, 1995E \citep{Quinn06}, 1998bu \citep{Garnavich01,Cappellaro01}, 2002hh \citep{Welch07,Otsuka12}, 2003gd \citep{Sugerman05,VanDyk06,Otsuka12}, 2004et \citep{Otsuka12}, 2006X \citep{Wang08,Crotts08}, 2006bc \citep{Gallagher11,Otsuka12}, 2006gy \citep{Miller10}, 2007it \citep{Andrews11}, and 2008bk \citep{VanDyk13}. All of the aforementioned LEs had the common selection criterion that they were found serendipitously while the transient source was still bright. Early on in the last century, \cite{Zwicky40} had the idea that it might be possible to learn more about historical SNe by studying their scattered LEs. However, the few dedicated surveys trying to implement this idea for historic SNe \citep{vandenBergh65b,vandenBergh65a,vandenBergh66,Boffi99} and novae \citep{vandenBergh77,Schaefer88} were not successful. With the emergence of CCDs as astronomical detectors in combination with the advancement in telescope technology that allowed to image larger field-of-views, the wide-field time-domain surveys at the beginning of this century significantly improved in depth and area. These improvements led to the first discoveries of LEs at angular distances from the ancient transients too large to suggest an immediate association. The i400-900 year-old LEs from three LMC supernovae were found by \cite{Rest05b} as part of the SuperMACHO survey \citep{Rest05a}. Subsequent targeted searches in our Galaxy found LEs of Tycho's SN \citep{Rest07,Rest08b}, Cas~A \citep{Rest07,Rest08b,Krause08a}, and $\eta$~Carinae \citep{Rest12_etac}. | In the last decade, LE spectroscopy has emerged as powerful tool to spectroscopically classify ancient and historic SNe, for which no contemporary observations with modern instrumentation were possible. More recently, the technique of LE spectroscopy has been refined and improved, and it is now possible to utilize it to directly probe the asymmetries of the transient sources of the LEs. Furthermore, a spectroscopic time series can be obtained for favorable scattering dust filament structure or if the source transient is an event with a long time-scale. | 14 | 3 | 1403.7223 |
1403 | 1403.6680_arXiv.txt | We present lightcurve observations and their multiband photometry for near-Earth object (NEO) 2011~$\mathrm{XA_3}$. The lightcurve has shown a periodicity of 0.0304 $\pm$ 0.0003 day (= 43.8 $\pm$ 0.4 min). The fast rotation shows that 2011~$\mathrm{XA_3}$ is in a state of tension (i.e. a monolithic asteroid) and cannot be held together by self-gravitation. Moreover, the multiband photometric analysis indicates that the taxonomic class of 2011~$\mathrm{XA_3}$ is S-complex, or V-type. Its estimated effective diameter is 225 $\pm$ 97~m (S-complex) and 166 $\pm$ 63~m (V-type), respectively. Therefore, 2011~$\mathrm{XA_3}$ is a candidate for the second-largest fast-rotating monolithic asteroid. Moreover, the orbital parameters of 2011~$\mathrm{XA_3}$ are apparently similar to those of NEO (3200) Phaethon but F/B-type. We computed the orbital evolutions of 2011~$\mathrm{XA_3}$ and Phaethon. However, the results of the computation and distinct taxonomy indicate that neither of the asteroids is of common origin. | The physical properties of asteroids provide us with important clues to clarify the compositions, strengths, and impact history of planetesimals that were formed in the early solar system. One representative methods for investigating the physical properties of asteroids is their lightcurve observations. Lightcurve observations are able to deduce the rotational status and shapes of asteroids. Past observations found that diameters of most asteroids rotating shorter than 2.2~h were smaller than 200~m \citep{Pravec00}. The fast-rotating asteroids have structurally significant tensile strength (so-called monolithic asteroids) because such asteroids need to overcome their own centrifugal force. On the other hand, the slow rotating asteroids larger than 200~m in diameter maintain themselves by their gravity or the cohesive force of bonded aggregates \citep{Ric09}. Some such asteroids are structurally rubble pile \citep{Abe06}. The near-Earth object (NEO), 2001~OE$_{84}$, has the largest diameter among fast-rotating asteroids whose rotational periods are reasonably determined in the Asteroid Lightcurve Database (LCDB) with quality code U$\ge$2 \citep{Warner09}. When the geometric albedo $p_{V}$ is 0.18, its diameter and rotational period are $\sim 700$~m and 29.1909~min, respectively \citep{Pravec02}. \cite{Hicks09} has reported that the diameter and rotational period of the NEO 2001~FE$_{90}$ are 200~m and 28.66~min, assuming the geometric albedo $p_{V}$ of 0.25. The diameter and rotational period of 2001~VF$_2$ are 1.39~h and 145--665~m \citep{Her11}. Despite searches having been devoted to their detection, few fast-rotating asteroids clearly larger than 200 m have been found \citep{Pravec00b, Whiteley02, Kwiat10a, Kwiat10b}. A numerical simulation indicates that fast-rotating asteroids are provided by the collisional disruption of parent bodies \citep{Asphaug99}. The relationship between the diameter and rotational period of ejecta in a laboratory impact experiment is similar to that of monolithic asteroids \citep{Kado09}. Moreover, explorations by NEAR and Hayabusa spacecrafts have discovered that the shapes of boulders on NEOs (433) Eros (mean diameter of 16.84~km) and (25143) Itokawa (mean diameter of 0.33~km) resembled those of monolithic asteroids \citep{Michi10}. Therefore, monolithic asteroids are thought to be generated by impact craterings or catastrophic disruptions on the parent bodies. The Rosetta and the NEAR spacecraft have determined the size frequency distributions of boulders on the main belt asteroid (MBA) (21) Lutetia (mean diameter of 95.76~km) and Eros \citep{Thomas01, Kup12}. In the case of Lutetia, the largest boulder is $\sim 300$~m in diameter. The power law index $\alpha$ of the cumulative size-frequency distribution, $N(>D)$ $\propto$ $D$$^{-\alpha}$, is about -5 and becomes shallower in regions smaller than 150~m, where $N$ is the number of boulders and $D$ is the diameter of boulders. A similar trend in slope is also shown for Eros. When we hypothesize that the size frequency distributions of the boulders on Lutetia and Eros are recognized as the typical size frequency distribution of impact ejecta, some ejecta smaller than 150~m can escape from the parent object that are likely to become small and fast-rotating asteroids. To confirm such a hypothesis, we need to show clearly how the fast-rotating asteroids inhabit from 150 m or less to subkilometer size by examining gradual but steady observational data accumulations. \\ The purpose of our study is to obtain the rotational period, the taxonomic class for NEO 2011~$\mathrm{XA_3}$. Since apparent brightness increases when NEOs closely encounter the Earth, the observations of NEOs are suitable for the elucidation of physical properties of subkilometer-sized asteroids. Revealing the taxonomic class confines the albedo, and increases the estimate accuracy for the diameter. 2011~$\mathrm{XA_3}$ was discovered by Pan-STARRS (the Panoramic Survey Telescope \& Rapid Response System) on December 15, 2011. The absolute magnitude $H$ = 20.402 $\pm$ 0.399 (JPL Small-Body database\footnote{http://ssd.jpl.nasa.gov/sbdb.cgi}) indicates that the object is a subkilometer-sized asteroid. Based on more than 100 astrometric observations for this asteroid, the orbital parameters are reduced to $a$ = 1.48 AU, $e$ = 0.93, $i$ = $28\fdg1$, $\Omega$ = $273\fdg6$ and $\omega$ = $323\fdg8$. Its orbit is apparently similar to those of NEO (3200) Phaethon, $a$ = 1.27 AU, $e$ = 0.89, $i$ = $22\fdg2$, $\Omega$ = $265\fdg3$ and $\omega$ = $322\fdg1$, though the semi-major axis has a difference. Phaethon is the parent body of the Geminids, one of the most intense meteor showers throughout the year. Moreover, the Phaethon Geminid Complex (hearafter, PGC) like 2005 UD and 1999 YC, which are objects dynamically linked with Phaethon and the Geminids, are proposed \citep{Ohtsuka06, Ohtsuka08}. In fact, B/F/C taxonomy of 2005 UD and 1999 YC is shared by Phaethon \citep{Je06,Kinoshita07, Ka08}. We also investigate whether 2011~$\mathrm{XA_3}$ is a member of PGC by the multiband photometry and by calculating the orbital motion of both 2011~$\mathrm{XA_3}$ and Phaethon. In this paper, we deal with the following: In Section 2, we describe the observations and their data reduction. In Section 3, we mention the results of rotational period, taxonomic class, and the orbital simulation. In Section 4, we discuss the relationship between 2011~$\mathrm{XA_3}$ and Phaethon. Moreover, we focus on the heating effect due to the close perihelion distance of 2011~$\mathrm{XA_3}$. Finally, we summarize the physical properties of 2011~$\mathrm{XA_3}$ and mention the observation efficiency of subkilometer-sized NEO using small- and medium-aperture telescopes. | \subsection{Relationship with Phaethon} As we mentioned in the section above, 2011~$\mathrm{XA_3}$ is not a PGC member based on the taxonomic analysis and the orbital calculation. Nevertheless, Phaethon with nominal orbital elements has more opportunities for close encounters with 2011~$\mathrm{XA_3}$ than with either the Earth or the Moon: 59 times within $0.05$ AU in the past $30,000$ years with 2011~$\mathrm{XA_3}$, in comparison with 23 times with the Earth and Moon. This reminds us of the possibility of a collisional event. In recent years, an unexpected brightening and a comet-like dust tail were detected on Phaethon around the perihelion \citep{Li13b, Je13}. This was because thermal fractures and decomposition cracking of hydrated minerals produced the dust ejections, whether the amount of dust production was enough to supply the Geminids in steady state or not was not concluded. The sublimation of ice inside the PGC precursor object due to the thermal evolution has been suggested as a possible mechanism for the breakup of PGC precursor object \citep{Ka09}. Alternatively, impacts event cannot be denied as a possible breakup mechanism. In addition, it is interesting that there exists a rotationally S-type-like color region on Phaethon's surface \citep{cochran84}. Therefore, we infer a $priori$ that 2011~$\mathrm{XA_3}$ might be a remnant candidate impacted with a potential PGC precursor. \subsection{Surface Material} Intense solar radiation heating by the perihelion distance of 0.11 AU elevates temperature of the surface material of 2011~$\mathrm{XA_3}$. Assuming 2011~$\mathrm{XA_3}$ as S-type asteroid $p_{V}$ = 0.209 and $G$ = 0.24, we can estimate the FRM (fast rotating model) surface temperature \citep{Leb89} heating up to 900~K. 2011 $\mathrm{XA_3}$ is heated repeatedly when it comes close to the Sun and the duration at 900~K reaches at least 3,000~yrs (Figure 4), which could result in unique mineralogy of the surface material of 2011~$\mathrm{XA_3}$. A recent sample return mission from S-type asteroid Itokawa revealed that continuous reduction reaction from FeO to metallic Fe due to solar proton implantation to silicates is a main mechanism for developing reduction rims of the silicate crystals on the regolith surface \citep{Noguchi11}. The reduction rims are responsible for space weathering of S-type asteroids. The same reactions are expected to occur on the surface of 2011~$\mathrm{XA_3}$ at a much higher weathering rate, because the asteroid is much closer to the Sun and the reaction takes place at an elevated temperature. At a low-temperature surface of silicates, the space weathering process makes the reduction rims, consisting of amorphous silicates and metallic iron, from FeO-bearing crystalline silicates \citep{Noguchi11} . On the other hand, at a high-temperature surface of silicates on the asteroid 2011 $\mathrm{XA_3}$, it is expected that the reduction rims thicken, the amorphous silicates crystallize, and the small Fe particles integrate. Our results show that 2011~$\mathrm{XA_3}$ shows a reflectance spectrum feature intermediate between S- and V-type (Figure 2). A meteoritic analog for the surface composition of S-type corresponds to ordinary chondrites that consist mainly of olivine, pyroxene, and plagioclase. The following reaction would take place if olivine were present on the surface: \begin{equation} {\rm (Mg,Fe)_2SiO_4} + {\rm H_2} \longrightarrow {\rm MgSiO_3} + {\rm Fe} +{\rm H_2O}. \end{equation} Since the reaction proceeds at high temperature, H$_2$O evaporates away from the surface and MgSiO$_3$ crystallizes to pyroxene for a short time. Therefore, pyroxene/olivine ratio and Fe metal abundance are expected to increase on the surface of 2011~$\mathrm{XA_3}$. Interestingly, the increase of pyroxene/olivine ratio is an opposite trend induced by early thermal metamorphism that occurred in the interior of an S-type asteroid 4.6 billion years ago \citep{Gas02}. As for other components of S-type asteroids, pyroxene is more difficult to be reduced than olivine \citep[e.g.,][]{Sin03} and plagioclase does not contain FeO to be reduced. Therefore a major change in mineralogy of the surface of ``S-type" 2011 $\mathrm{XA_3}$ is an increase in the pyroxene/olivine ratio and Fe-metal abundance. If 2011~$\mathrm{XA_3}$ is a V-type asteroid whose meteorite analog is HED (Howardrites, Eucrites, and Diogenites) meteorites, consisting mainly of pyroxene and plagioclase, in this case the mineralogical change is limited to further crystallization of individual minerals. The increase in the pyroxene/olivine ratio during high-temperature space weathering of S-type asteroids makes the mineral assemblage similar to V-type asteroids, which are rich in pyroxene. 2011~$\mathrm{XA_3}$ may be an S-type asteroid with a high pyroxene/olivine or a V-type asteroid. In either case, the reflectance spectrum is similar and difficult to distinguish as we observed in the color-color diagram (Figure 2). Last, we mention the possibility of Q-type because the color-color diagram of 2011~$\mathrm{XA_3}$ indicates the intermediate between S- and V-type, and the population ration of Q-type asteroids is dominant in the NEO region than R- and O-type. The cause of 2011~$\mathrm{XA_3}$ being a monolith is thought to the rotational fission of a rubble pile object due to the YORP effect, and the ejector by impact craterings or catastrophic disruption on the parent bodies. As we describe above, the heating to 900~K promotes the space weathering. However, if the rotational fission and the ejection by impacts took place recently, the surface of 2011~$\mathrm{XA_3}$ has not been long exposed to the solar radiation. In that case, the surface color of 2011~$\mathrm{XA_3}$ might indicate Q-type. | 14 | 3 | 1403.6680 |
1403 | 1403.6964_arXiv.txt | The Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) has observed more than 80,000 solar energetic events since its launch on February $12^{\rm th}$, 2002. Using this large sample of observed flares, we studied the spatio-temporal relationship between succeeding flares. Our results show that the statistical relationship between the temporal and spatial differences of succeeding flares can be described as a power law of the form $R(t)\sim t^p$ with $p=0.327\pm 0.007$. We discuss the possible interpretations of this result as a characteristic function of a supposed underlying physics. Different scenarios are considered to explain this relation, including the case where the connectivity between succeeding events is realised through a shock wave in the post Sedov-Taylor phase or where the spatial and temporal relationship between flares is supposed to be provided by an expanding flare area in the sub-diffusive regime. Furthermore, we cannot exclude the possibility that the physical process behind the statistical relationship is the reordering of the magnetic field by the flare or it is due to some unknown processes. | Sudden or explosive releases of energy are common phenomena in different cosmic objects occurring within wide range of energy release rate from as low as $10^{17}$ J sec$^{-1}$ in case of the Sun until as high as $10^{47}$ J sec$^{-1}$ in gamma ray bursts at cosmological distances. In the present context, the notion of sudden or explosive event means that the ratio of the released energy ($E$) compared to rate of its change ($dE/dt$) is significantly shorter that the dynamical time scale of the object, i.e. the ratio of the characteristic physical size ($R$) and the propagation speed ($v$) of any kind of disturbance initiated during the release (see, e.g. \citealt{b38}). A typical characteristic of these explosions is that the rising time of energy release is very short in comparison with the decay phase independently of the physics behind these explosions. Our study will focus on sudden energy releases in the solar atmosphere called flares. The systematic multi-wavelength study of flares revealed many details that help us in understanding the physics behind these phenomena, however, there are still many unanswered questions concerning the true dynamics and energetics of flares (e.g. \citealt{b1, SM11, Fea11}). It is widely accepted that magnetohydrodynamic (MHD) processes are responsible for producing a flare. The energy released in a flare could originate from magnetic reconnection at the top of magnetic loop where the magnetic flux tubes take an X-type configuration. During reconnection events, the magnetic energy stored in magnetic field lines is released in a very localised way into thermal and kinetic energy. This energy is transported from the site of reconnection via, e.g. radiation, slow and fast MHD shock waves, accelerated particles, and high-speed collimated hot plasma flows (jets). Flares usually occur in the solar corona but part of the released energy is transported downward to the lower atmosphere where large-scale disturbances were observed travelling in a large distance away from the flaring sites as e.g. sunquakes observed in the photosphere \citep{ko98, b10} or Moreton waves in the chromosphere \citep{mo60}. The energy of flares can be within a range of $10^{17}-10^{26} \, J$ (\citealt{b7}). Observations often reveal pairs of flares that occur with close temporal and/or spatial proximity and a physical connection between events was supposed to exist. Flares that occur because of common physical reason in different active regions are called sympathetic flares, and large-scale coronal structures are the most probable causes of their connection (e.g. \citealt {Moon02, b55}). Recently \citet {Liuea09} found that four flares and two fast coronal mass ejections (CMEs) occurred with a causal relationship in an active region within $\sim$1.5 hr time interval. They called these types of solar flares occurring in the same active region with a causal relationship successive flares. \citet{Zuc09} also found a sequence of successive destabilisation of the magnetic field configuration starting with a filament eruption (relatively cool, dense object of chromospheric material suspended in the corona by magnetic fields) and ended in a large flare within $\sim$2 hrs; they referred to the process as domino effect. \citet{Jiang09} presented the evidence for occurrences of magnetic interactions between a jet, a filament and coronal loops during a complex event, in which two flares sequentially occurred at different positions of the same active region stating with 1 hr time difference and two associated CMEs. Recently \citet{b52} presented a study of a multiple flare activity containing three small flares, and a major eruptive flare over the period of two hours. They concluded that the small precursor flares (preflares) indicated the localised magnetic reconnections associated with different evolutionary stages of the filament in the pre-eruption phase and these events play a crucial role in destabilising the filament leading to a large-scale eruption. The preflare activity occurs in the form of discrete, localised X-ray brightenings observed between 2 and 50 min before the impulsive phase of the flare and filament acceleration. \citet{Chifor07} claim that the X-ray precursors provide evidence for a tether-cutting mechanism initially manifested as localised magnetic reconnection being a common trigger for both flare emission and filament eruption. \citet{Kim08} also demonstrated that a preflare eruption and the main flare have a causal relation because they are triggered by a sequential tether-cutting process. Based on these observations and conclusions, it is natural to question whether there are a significant number of flares physically connected in such way that their relationship can be revealed with statistical methods. This question is extensively debated in the literature but the statistical results presented so far have left this question open. The used methods usually focus on the study of flare waiting times distribution (WTD, the distribution of times between events) which can provide information about whether flares are independent events, or not (e.g. \citealt {W00, Moon02}). The results suggests that determination of WTD gives varied results, suggesting that the observed distribution may depend on the particular active region, on time, and that it also may be influenced by event definition and selection procedures \citep {W09}. The solar flare sympathy is probably a statistically weak effect \citep {WC06}, but the successive flares probably do not occur randomly in time and the WTD are regulated by solar flare mechanisms \citep {Kubo08}. We apply an alternative statistical method that takes into account both the spatial and temporal distributions of flares. The observational data used in the present study are the flares appearing in the list provided by the RHESSI satellite between 2002 and 2010. The paper is structured as follows: in Section \ref{redat} we present the observational set-up of our study and data used in the analysis. Section \ref{srel} deals with the statistical relationship between spatial and temporal differences of succeeding flares. The discussion on the possible physical interpretations of the statistical results is given in Section \ref{disc}. Section \ref{sum} presents the summary of the main results and the conclusions. The paper includes an Appendix containing a gallery of scatter plots for different time intervals displaying the statistical relationship discussed in the paper. | \label{sum} In the present paper we studied the statistical relationship between succeeding flare events recorded by the RHESSI satellite. We selected the flares that were associated to active regions in the period 2002-2010. The heliographic coordinates of events were calculated applying the appropriate corrections. With the help of the coordinates we were able to determine spatial distances between succeeding events ($ds$) in the same active regions, while temporal differences ($dT$) were deducted from the RHESSI catalog. We studied the statistical relationship between the $ds$ and $dT$ differences dividing the RHESSI data in subsamples according to the phase of the solar cycle and the energy channels, respectively. If we take into account only flares occurring in the same active region, we obtain a linear relationship in the $\{\log_{10}(ds)-\log_{10}(dT)\}$ diagram indicating a power law connection between the spatial and temporal differences of succeeding flares. This relationship may reveal a hidden factor responsible for the statistical connection. Performing a Principal Component Analysis (PCA) on the $ds$ and $dT$ values we obtained the parameters of the power law relationship. The parameters obtained in this way showed a remarkable homogeneity, independently from the phase of the solar cycle. Several possibilities were discussed to explain the temporal and spatial correlation between events. One candidate that could explain the connectivity was a blast wave. This assumption was made based on the statistically determined time-distance relationship of the form $R(t)=At^p$, where $p=0.327\pm 0.007$, i.e. very close to the typical value of a shock in the post Sedov-Taylor stage. The most problematic point of this scenario is the derived propagation speed of these shocks being sub-sonic, at least a few orders of magnitude less than expected. We had to assume that the blast wave transmitting the causal connection between two succeeding flares takes only a few percent of their spatial and temporal differences if it plays a role at all. It seems more possible, that the net effect of different, basically magnetic, processes in the solar atmosphere mimics statistically a functional relationship between the time and spatial distances of succeeding flares in a mathematical form of a blast wave. In particular, the relationship between temporal and spatial distances we observed could be the characteristics of the magnetic field re-arrangement after a flare in the same active region. Finally, we need to note that our derived values are similar to the values obtained earlier by \citet {a12} using SDO data, where the spatial and temporal differences were connected to the expansion of the active region area in the sub-diffusive stage. If our data analysis would point to the same mechanism described by this author, it would be a further evidence for the universal character of it as, according to our findings, this occurs not only in particular active regions but it can be shown to be a trend over much longer periods. | 14 | 3 | 1403.6964 |
1403 | 1403.4961_arXiv.txt | We present the redshift evolutions and distributions of the gamma-ray luminosity and photon spectral index of flat spectrum radio quasar (FSRQ) type blazars, using non-parametric methods to obtain the evolutions and distributions directly from the data. The sample we use for analysis consists of almost all FSRQs observed with a greater than approximately 7$\sigma$ detection threshold in the first year catalog of the {\it Fermi} Gamma-ray Space Telescope's Large Area Telescope, with redshfits as determined from optical spectroscopy by Shaw et al. We find that FSQRs undergo rapid gamma-ray luminosity evolution, but negligible photon index evolution, with redshift. With these evolutions accounted for we determine the density evolution and luminosity function of FSRQs, and calculate their total contribution to the extragalactic gamma-ray background radiation, resolved and unresolved, which is found to be 16(+10/-4)\%, in agreement with previous studies. | \label{intro} The majority of the extragalactic sources observed by the Large Area Telescope (LAT) on the {\it Fermi} Gamma-ray Space Telescope are blazars \citep[e.g.][]{Fermiyr1}, the type of active galactic nuclei (AGNs) in which one of the jets is aligned with our line of sight \citep[e.g.][]{BK79}. Among AGN only blazars frequently feature prominent gamma-ray emission, and the gamma-ray emission is an essential observational tool for contstraining the physics of the central engines of AGNs \citep[e.g.][]{Dermer07}. Understanding the characteristics of blazars is also crucial for evaluating their contribution as a source class to the extragalactic gamma-ray background (EGB) radiation. In \citet{BP1} --- hereafter BP1 --- we explored the source counts (the so-called Log$N-$Log$S$ relation) of {\it Fermi}-LAT blazars, using those detected with a greater than approximately seven sigma detection threshold in the first-year {\it Fermi}-LAT catalog, and determined their contribution to the EGB. Estimating the total contribution of blazars to the EGB required an extrapolation of the source counts to lower fluxes, below those detected by the {\it Fermi}-LAT. However, in the presence of luminosity and/or density evolution with redshift, a more accurate estimate of the integrated flux from blazars requires determining and factoring in the evolution of blazars with redshift. \citet{Shaw} provide spectroscopically-determined redshifts for almost all of the FSRQ blazars from the {\it Fermi}-LAT first year catalog. With the inclusion of these redshifts, we can apply our techniques to determine the evolutions of the luminosity and photon index with redshift, the density evolution, and the distributions of luminosity and photon index for FSRQs. Fluxes for {\it Fermi}-LAT sources are measured and reported for a given photon energy range. The lower limit flux for detection of blazars by the {\it Fermi}-LAT depends strongly on a source's gamma-ray spectrum, such that objects with harder spectra are able to be detected above the background level at lower fluxes than those with softer spectra \citep{Atwood09}. This means that for determination of the luminosity distribution one needs both a measure of the flux and the photon index $\Gamma$, and that then one deals with a bi-variate distribution of fluxes and indexes, which is truncated because of the this observational bias in the flux-index plane (as seen in Figure \ref{lumsandsis}), often referred to as Malmquist bias. Additionally, of course, there is a truncation in the luminosity-redshift plane arising because of the relationship between flux and luminosity. Obtaining a bias free determination of the distributions of luminosity and photon index is therefore necessarily quite a bit more complicated than a simple counting of sources. \citet{Marco} use simulated data to account for the detection biases in analyzing this data. Here we use non-parametric methods to determine the luminosity and density distributions directly from the observational data. When dealing with a multivariate distribution, the first required step is the determination of the correlation (or statistical dependence) between the variables, which cannot be done by simple procedures when the data is truncated \citep[e.g.][]{P92}. We use the procedures developed by Efron and Petrosian \citep[EP,][]{EP92,EP99} and extended by \citet{QP1,BP1,QP2} to account reliably for the complex observational selection biases to determine first the intrinsic correlations (if any) between the variables. These techniques have been proven useful for application to many sources with varied characteristics, including to the Log$N-$Log$S$ relation for blazars in BP1, and to radio and optical luminosity in quasars in \citet{QP1} and \citet{QP2}, where references to earlier works are presented. In this paper we apply these methods to determine the luminosity and photon index evolutions of {\it Fermi}-LAT blazars, as well as the density evolution, and local ($z$=0) gamma-ray luminosity function. In \S \ref{datasec} we discuss the data used, and in \S \ref{evs} we explain the techniques used and present the results for the luminosity and photon index evolution. In \S \ref{dev} and \S \ref{ll} we present the density evolution and local luminosity function and photon index distribution. In \S \ref{bgnd} we calculate the total contribution of FSRQs to the extragalactic gamma-ray background (EGB) radiation. This paper assumes the standard cosmology throughout. | \label{disc} We have used a rigorous method to calculate non-parametrically and directly from the data the redshift evolutions of the gamma-ray luminosity and photon index, as well as the density evolution, gamma-ray LF, and contribution to the EGB of FSRQ blazars. We use a data set consisting of the FSRQs in the {\it Fermi}-LAT first year extragalactic source catalog with $TS \geq 50$ and which lie at Galactic latitude $\vert b \vert \geq 20^{\circ}$, with spectroscopically determined redshifts provided by \citet{Shaw}. The method employed accounts robustly for the pronounced data truncation introduced by the selection biases inherent in the {\it Fermi}-LAT observational catalog. The reliability of the methods employed has been demonstrated and discussed in \citet{QP1}, \citet{QP2}, and the appendix of BP1. We note that since spectroscopic redshifts are available for virtually all of the FSRQs in the {\it Fermi}-LAT first year extragalactic source catalog, there is not a relevant limiting optical flux which must be considered, and the only relevant truncation in the data is that arising from gamma-ray flux and photon index. In \S 3.3 of BP1 we discuss the sources of error that may affect these determinations, including measurement uncertainties, blazar variability, and source confusion. As discussed there, in the determination of the contribution of FSRQs to the EGB, these will be sub-dominant to the uncertainty resulting from the range of mean values of the photon index distribution that we consider, as well as the high redshift end of the density evolution function. Subtler issues affecting the derived distributions may arise because of the finite bandwidth of the {\it Fermi}-LAT and lack of complete knowledge of the objects' spectra over a large energy range and deviations from simple power laws. This has the greatest potential to effect determinations of the photon index distribution, which we do not determine in this work but carry over from BP1. However the bandwidth 100 MeV to 100 GeV is wide enough that the contribution of sources which peak outside of this range to the LF and evolution, density evolution, and the EGB in this energy range will be small. We find that FSRQs have a strong luminosity evolution with redshift, well characterized (at low redshifts) by the evolution factor (1+$z$)$^{k_L}$ with $k_L$=5.5$\pm$0.5. This, along with positive evolutions in other wavebands such as optical and radio \citep[e.g.][]{QP2} and X-ray \citep[e.g.][]{Aird10}, favors models in which at higher redshifts AGN systems featured on average more massive black hole and accretion disk systems, and/or more rapidly rotating black holes. We find that FSRQs do not exhibit appreciable photon index evolution with redshift, indicating that the mean spectrum of accelerated high energy particles from AGN central engines has remained constant over the history of the Universe. Given these evolutions, the density evolution, and the local LF and photon index distribution, we determine that the total energy density from FSRQs at all redshifts and luminosities is $\mathcal{I}_{\rm \gamma : FSRQs}$=7.6 (+4.0/-1.1) $\times$10$^{-4}$ MeV cm$^{-2}$ sec$^{-1}$ sr$^{-1}$, which is 16 (+10/-4)\% of the total EGB. This indicates that FSRQ blazars are a significant, but not dominant, component of the EGB. | 14 | 3 | 1403.4961 |
1403 | 1403.6155_arXiv.txt | We extend the non-equilibrium model for the chemical and thermal evolution of diffuse interstellar gas presented in \citet{richings14} to account for shielding from the UV radiation field. We attenuate the photochemical rates by dust and by gas, including absorption by H\textsc{i}, H$_{2}$, He\textsc{i}, He\textsc{ii} and CO where appropriate. We then use this model to investigate the dominant cooling and heating processes in interstellar gas as it becomes shielded from the UV radiation. We consider a one-dimensional plane-parallel slab of gas irradiated by the interstellar radiation field, either at constant density and temperature or in thermal and pressure equilibrium. The dominant thermal processes tend to form three distinct regions in the clouds. At low column densities cooling is dominated by ionised metals such as Si\textsc{ii}, Fe\textsc{ii}, Fe\textsc{iii} and C\textsc{ii}, which are balanced by photoheating, primarily from H\textsc{i}. Once the hydrogen-ionising radiation becomes attenuated by neutral hydrogen, photoelectric dust heating dominates, while C\textsc{ii} becomes dominant for cooling. Finally, dust shielding triggers the formation of CO and suppresses photoelectric heating. The dominant coolants in this fully shielded region are H$_{2}$ and CO. The column density of the H\textsc{i}-H$_{2}$ transition predicted by our model is lower at higher density (or at higher pressure for gas clouds in pressure equilibrium) and at higher metallicity, in agreement with previous PDR models. We also compare the H\textsc{i}-H$_{2}$ transition in our model to two prescriptions for molecular hydrogen formation that have been implemented in hydrodynamic simulations. | The thermal evolution of gas is an important component of hydrodynamic simulations of galaxy formation as it determines how quickly the gas can cool and collapse to form dense structures, and ultimately stars. The star formation in such simulations can be limited to the cold phase of the interstellar medium (ISM) if we have sufficient resolution to resolve the Jeans mass ($M_{\rm{J}} \propto \rho^{-1/2} T^{3/2}$) in the cold gas. It is desirable to include a multi-phase treatment of the ISM, as this will produce a more realistic description of the distribution of star formation within the galaxy, along with the resulting impact of stellar feedback on the ISM and the galaxy as a whole \citep[e.g.][]{ceverino09,governato10,halle13,hopkins13}. Therefore, it is important that we correctly follow the thermal evolution of gas between the warm ($T \sim 10^{4}$ K) and cold ($T \la 100$ K) phases of the ISM in these simulations. The radiative cooling rate depends on the chemical abundances in the gas, including the ionisation balance, and hence on its chemical evolution. However, following the full non-equilibrium chemistry of the ISM within a hydrodynamic simulation can be computationally expensive, as it requires us to integrate a system of stiff differential equations that involves hundreds of species and thousands of reactions. Therefore, many existing cosmological hydrodynamic simulations use tabulated cooling rates assuming chemical (including ionisation) equilibrium. For example, the cosmological simulations that were run as part of the OverWhelmingly Large Simulations project \citep[OWLS;][]{schaye10} use the pre-computed cooling functions of \citet{wiersma09}, which were calculated using \textsc{Cloudy}\footnote{\url{http://nublado.org/}} \citep{ferland98,ferland13} as a function of temperature, density and abundances of individual elements assuming ionisation equilibrium in the presence of the \citet{haardt01} extragalactic UV background. However, this assumption of ionisation equilibrium may not remain valid if the cooling or dynamical time-scale becomes short compared to the chemical time-scale \citep[e.g.][]{kafatos73,gnat07,oppenheimer13a,vasiliev13} or if the UV radiation field is varying in time \citep[e.g.][]{oppenheimer13b}. In the first paper of this series (\citealt{richings14}; hereafter paper I) we presented a chemical network to follow the non-equilibrium thermal and chemical evolution of interstellar gas. Using this model we investigated the chemistry and cooling properties of optically thin interstellar gas in the ISM and identified the dominant coolants for gas exposed to various UV radiation fields. We also looked at the impact that non-equilibrium chemistry can have on the cooling rates and chemical abundances of such gas. In this paper we extend our thermo-chemical model to account for gas that is shielded from the incident UV radiation field by some known column density. We focus on physical conditions with densities $10^{-2} \, \text{cm}^{-3} \la n_{\rm{H_{tot}}} \la 10^{4} \, \text{cm}^{-3}$ and temperatures of $10^{2} \, \text{K} \la T \la 10^{4} \, \text{K}$. This is most relevant to gas that is cooling from the warm phase to the cold phase of the ISM. We apply our model to a one-dimensional plane-parallel slab of gas that is irradiated by the \citet{black87} interstellar radiation field to investigate how the chemistry and cooling properties of the gas change as it becomes shielded from the UV radiation, both by dust and by the gas itself. The spectral shape of the radiation field will change with the depth into the cloud as high energy photons are able to penetrate deeper. Hence, the major coolants and heating processes will vary with column density. Such one-dimensional models are commonly used to model photodissociation regions (PDRs) and diffuse and dense clouds \citep[e.g.][]{tielens85,vandishoeck86,vandishoeck88,lepetit06,visser09,wolfire10}. It has been suggested recently that the star formation rate of galaxies may be more strongly correlated to the molecular gas content than to atomic hydrogen \citep{wong02,schaye04,kennicutt07,leroy08,bigiel08,bigiel10}, although the more fundamental and physically relevant correlation may be with the cold gas content \citep{schaye04,krumholz11a,glover12}. Motivated by this link between molecular hydrogen and star formation, a number of studies have implemented simple methods to follow the abundance of H$_{2}$ in numerical simulations of galaxies \citep[e.g.][]{pelupessy06,gnedin09,mckee10,christensen12}. We compare the H$_{2}$ fractions predicted by some of these methods to those calculated using our model to investigate the physical processes that determine the H\textsc{i}-H$_{2}$ transition and to explore in which physical regimes these various prescriptions remain valid. This paper is organised as follows. In section~\ref{modelSummary} we summarise the thermo-chemical model presented in paper I, and in section~\ref{shielding_section} we describe how the photochemical rates are attenuated by dust and gas. We look at the photoionisation rates in section~\ref{shieldIon_section}, the photodissociation of molecular species in section~\ref{shieldDissoc_section} and the photoheating rates in section~\ref{photoheat_section}. In section~\ref{shieldResultsSect} we apply this model to a one-dimensional plane-parallel slab of gas to investigate the chemistry and cooling properties of the gas as it becomes shielded from the UV radiation field, and we compare these results with \textsc{Cloudy}. In section~\ref{H2modelsSection} we compare the time-dependent molecular H$_{2}$ fractions predicted by our model with two prescriptions for H$_{2}$ formation taken from the literature that have been implemented in hydrodynamic simulations. Finally, we discuss our results and conclusions in section~\ref{conclusions}. | We have extended the thermo-chemical model from paper I to account for gas that becomes shielded from the incident UV radiation field. We attenuate the photoionisation, photodissociation and photoheating rates by dust and by the gas itself, including absorption by H\textsc{i}, H$_{2}$, He\textsc{i}, He\textsc{ii} and CO where appropriate. For the self-shielding of H$_{2}$, we use a new temperature-dependent analytic approximation that we fit to the suppression of the H$_{2}$ photodissociation rate predicted by \textsc{Cloudy} as a function of H$_{2}$ column density (see appendix~\ref{H2self_comparison_section}). Using this model, we investigated the impact that shielding of both the photoionising and the photodissociating radiation has on the chemistry and the cooling properties of the gas. We have performed a series of one-dimensional calculations of a plane-parallel slab of gas illuminated by the \citet{black87} interstellar radiation field at constant density. Comparing equilibrium abundances and cooling and heating rates as a function of column density with \textsc{Cloudy}, we generally find good agreement. At $n_{\rm{H_{tot}}} = 100 \, \text{cm}^{-3}$, solar metallicity and a constant temperature $T = 300$ K, we find that the H\textsc{i}-H$_{2}$ transition occurs at a somewhat lower column density in our model than in \textsc{Cloudy}'s big H2 model, with a molecular hydrogen fraction $x_{\rm{H_{2}}} = 0.5$ at $N_{\rm{H_{tot}}} \approx 8.1 \times 10^{19} \, \rm{cm}^{-2}$ in our model, compared to $N_{\rm{H_{tot}}} \approx 2.8 \times 10^{20} \, \rm{cm}^{-2}$ in \textsc{Cloudy} (see figure~\ref{shieldedCombinedFig}). However, \textsc{Cloudy}'s small H2 model predicts an H\textsc{i}-H$_{2}$ transition column density that is closer to our value, with $x_{\rm{H_{2}}} = 0.5$ at $N_{\rm{H_{tot}}} \approx 1.1 \times 10^{20} \, \rm{cm}^{-2}$. In the examples shown here, the H\textsc{i}-H$_{2}$ transition is determined by H$_{2}$ self-shielding in our model, as the residual molecular hydrogen fraction in the photodissociated region at low column densities is sufficient for self-shielding to become important before dust shielding. As the photodissociation rate is slightly lower in our model than in the big H2 model of \textsc{Cloudy}, the H$_{2}$ fraction at low column densities is slightly higher in our model. This explains why the transition column density is somewhat lower in our model compared to the \textsc{Cloudy} big H2 model. The importance of H$_{2}$ self-shielding also means that the molecular hydrogen transition is sensitive to turbulence, which can suppress self-shielding. The transition for carbon to form CO is primarily triggered by dust shielding and thus occurs at a higher column density, $N_{\rm{H_{tot}}} \sim 10^{22} \, \text{cm}^{-2}$. We also consider gas clouds with temperature and density profiles that are in thermal and pressure equilibrium, which is more realistic for a two-phase ISM than a gas cloud with constant temperature and density throughout. The effect of H$_{2}$ self-shielding is weaker in these examples due to the lower densities in the photodissociated region (see figure~\ref{ShieldIsobaricFig}). However, the H\textsc{i}-H$_{2}$ transition is still determined by self-shielding, which becomes more important in runs with higher pressure (and hence higher densities). This trend with density is consistent with previous studies that have looked at the importance of H$_{2}$ self-shielding for the H\textsc{i}-H$_{2}$ transition\citep[e.g.][]{black87b,draine96,lee96,krumholz09}. The H\textsc{i}-H$_{2}$ transition occurs at a lower total column density for higher density (or equivalently, for higher pressure if the cloud is in pressure equilibrium) and for higher metallicity (see figure~\ref{H2modelsFig}), in agreement with previous models of photodissociation regions \citep[e.g.][]{wolfire10}. These trends are due to the H$_{2}$ self-shielding, because an increase in the density and/or metallicity will increase the H$_{2}$ fraction in the photodissociated region, and hence decrease the total column density at which the H$_{2}$ becomes self-shielded. The time evolution of the H$_{2}$ fraction is also dependent on the density (or pressure) and the metallicity. For a gas cloud with pressure $P / k_{B} = 10^{3} \, \text{cm}^{-3} \, \text{K}$ and metallicity $0.1 Z_{\odot}$, the molecular hydrogen abundance in the fully shielded region only reaches equilibrium (starting from neutral, atomic initial conditions) after $\sim 1$ Gyr. This time-scale decreases as the pressure and/or the metallicity increase. We compare the dominant cooling and heating processes in our low pressure example ($P / k_{B} = 10^{3} \, \text{cm}^{-3} \, \text{K}$) at solar metallicity, and we find that they form three distinct regions (see figure~\ref{ShieldIsobaricFig}). At low column densities, where the dissociating and ionising radiation flux is still high, cooling is primarily from ionised metals such as Si\textsc{ii}, Fe\textsc{ii}, Fe\textsc{iii} and C\textsc{ii}, which are balanced by photoheating, primarily from H\textsc{i}. At column densities above $N_{\rm{H_{tot}}} \sim 2 \times 10^{20} \, \text{cm}^{-2}$ the hydrogen-ionising radiation above 1 Ryd becomes significantly attenuated by neutral hydrogen. This reduces the photoheating rates, making photoelectric dust heating the dominant heating mechanism, while C\textsc{ii} starts to dominate the cooling rate above $N_{\rm{H_{tot}}} \sim 10^{21} \, \text{cm}^{-2}$. It is also in this region that hydrogen becomes fully molecular, driven initially by an increase in the H\textsc{i} abundance and the rising density, and ultimately by self-shielding. Finally, dust shielding attenuates the radiation flux below 1 Ryd at column densities above $N_{\rm{H_{tot}}} \sim 10^{21} \, \text{cm}^{-2}$, which strongly cuts off the dust heating rate and also allows CO to form. In this fully shielded region heating is primarily from cosmic rays, while cooling is mostly from CO and H$_{2}$. Finally, we compare the H\textsc{i}-H$_{2}$ transition predicted by our one-dimensional plane-parallel slab simulations in thermal and pressure equilibrium with two other prescriptions for molecular hydrogen formation that are already employed in hydrodynamic simulations: the simpler non-equilibrium model of \citet{gnedin09} \textit{(Gnedin09)}, and the analytic equilibrium model developed in \citet{krumholz08,krumholz09} and \citet{mckee10} \textit{(KMT)} (see figure~\ref{H2modelsFig}). At low pressure ($P / k_{\rm{B}} = 10^{3} \, \rm{cm}^{-3} \, \rm{K}$) the equilibrium H$_{2}$ fractions predicted by all three models generally agree well, as does the time evolution of the H$_{2}$ fraction predicted by our model and the Gnedin09 model. However, at low metallicity (0.1 $Z_{\odot}$) cosmic ray dissociation of H$_{2}$ reduces the H$_{2}$ fraction in the fully shielded region by a factor of two in our model, but cosmic ray dissociation is not included in the KMT or the Gnedin09 models. At high pressure ($P / k_{\rm{B}} = 10^{5} \, \rm{cm}^{-3} \, \rm{K}$) the H\textsc{i}-H$_{2}$ transition predicted by our model in chemical equilibrium occurs at a slightly lower column density than in the KMT model (e.g. $f_{\rm{H_{2}}} = 0.5$ at $N_{\rm{H_{tot}}} = 2.2 \times 10^{19} \, \rm{cm}^{-2}$ at solar metallicity, compared to $N_{\rm{H_{tot}}} = 3.0 \times 10^{19} \, \rm{cm}^{-2}$ in the KMT model). Furthermore, the H\textsc{i}-H$_{2}$ transition at this high pressure is flatter in the Gnedin09 model than in our model or the KMT model, due to the different H$_{2}$ self-shielding function that they use. | 14 | 3 | 1403.6155 |
1403 | 1403.3400_arXiv.txt | Supernova (SN) 2008D/XRT 080109 is considered to be the only direct detection of a shock breakout from a regular SN to date. While a breakout interpretation was favored by several papers, inconsistencies remain between the observations and current SN shock breakout theory. Most notably, the duration of the luminous X-ray pulse is considerably longer than expected for a spherical breakout through the surface of a type Ibc SN progenitor, and the X-ray radiation features, mainly its flat spectrum and its luminosity evolution, are enigmatic. We apply a recently developed theoretical model for the observed radiation from a Wolf-Rayet SN exploding through a thick wind and show that it naturally explains all the observed features of SN 2008D X-ray emission, including the energetics, the spectrum and the detailed luminosity evolution. We find that the inferred progenitor and SN parameters are typical for an exploding Wolf-Rayet. A comparison of the wind density found at the breakout radius to the density at much larger radii, as inferred by late radio observations, suggests an enhanced mass loss rate taking effect about ten days or less prior to the SN explosion. This finding joins accumulating evidence for a possible late phase in the stellar evolution of massive stars, involving vigorous mass loss a short time before the SN explosion. | \label{sec:intro} The X-ray transient XRT 080109, associated with SN 2008D, was discovered by \emph{Swift}/XRT \citep{Soderberg+2008}. Later observations in Optical/UV led to its classification as a type Ibc SN, favoring a WR progenitor \citep{Soderberg+2008,Mazzali+2008,Malesani+2009,Modjaz+2009}. The X-ray signal of SN 2008D is the most convincing candidate for a shock breakout of a standard SN. Unlike other observed X-ray SN birth signals, which are all of rare broad line Ic SNe associated with Gamma-ray bursts \citep[e.g., SN2006aj/GRB060218;][]{Campana06}, the serendipitous discovery and the spectroscopic classification of SN 2008D suggest that it was a common signal, produced by a common Ibc SN \citep{Soderberg+2008}. Nevertheless, the detailed X-ray observations are still unexplained. The initial rise-time, the following light curve and the spectrum are different than the ones predicted for a standard SN breakout, namely a spherical breakout from the stellar surface \citep[e.g.][]{Chevalier+2008,Katz+2010,Nakar_Sari2010,Sapir+2011}. In order to explain the prolonged rise-time, both a breakout through a thick wind \citep{Soderberg+2008} and an aspherical breakout \citep{Couch+2011} were suggested. However, lacking detailed theoretical models of the two scenarios (a major progress in modeling aspherical breakouts was only recently achieved, by \citealt{Matzner+2013}), neither could be confronted with the detailed observed spectrum and light curve. As often happens in the absence of a consensual explanation, a burst of a mildly relativistic jet was also suggested \citep{Xu+2008,Li+2008,Mazzali+2008}, although this model has no predictions in terms of the expected X-ray emission, that can be compared to the observations. Recently we developed a detailed theoretical model for the emission from a SN breakout through a thick wind \citep{Svirski+2014}. Here we show, based on this model, that a scenario of a WR exploding through a thick wind naturally solves all inconsistencies, including the prolonged rise-time, and provides, without a need to invoke a significant breakout asphericity or an unconventional explosion scenario, an optimal explanation for the X-ray observations of SN 2008D. Moreover, the data fit all the model predictions, although these are tightly over-constrained, providing a strong support for this explanation. In Section \ref{sec:observations} we describe the observations and indicate their tension with a standard SN breakout interpretation. We then summarize, in Section \ref{sec:model}, our theoretical model for WR SNe exploding through a thick wind \citep{Svirski+2012,Svirski+2014} and show, in Section \ref{sec:impl}, that this scenario explains the X-ray observations. We conclude in Section \ref{sec:conclusion}. | \label{sec:conclusion} We offer a first coherent picture of SN 2008D X-ray observations, including a first explanation for the observed flat X-ray spectrum across the \emph{Swift}/XRT $0.3-10$ keV window during the first few hundred seconds, and a first account for the complete X-ray luminosity evolution of SN 2008D/XRT 080109, from breakout to $t\sim 1$ d. We find that a typical WR progenitor and typical SN parameters offer the optimal explanation for the observations. We apply our model for WR SNe exploding through a thick wind \citep{Svirski+2014} to the X-ray observations of SN 2008D. The model allows only little freedom in deriving physical parameters from the observations, and it over-constrains the shock velocity. In addition, it enforces a flat X-ray spectrum at $\tau \gtrsim 1$ and a $t^{-3}$ to $t^{-4}$ X-ray luminosity decay at $\tau\ll 1$. Hence, the combination of an observed $t_{bo}$ and an observationally inferred $v_{bo}$ uniquely determines the X-ray luminosity and spectrum as a function of time. Despite these tight constraints, the observations satisfy all the predictions of the model. Accordingly, SN 2008D had a compact progenitor ($R_*\lesssim3\times10^{11}$ cm), presumably a WR, that exploded through a thick wind of a standard wind density profile, $\rho\propto r^{-2}$, with $R_{bo}\approx 6\times10^{11}$ cm, $v_{bo}\approx 8\times 10^9 \,\rm{cm\,s^{-1}}$, $\tau_{bo}\approx 4$ and $n_{bo}\sim 10^{13}\,\rm{cm^{-3}}$. Observations of SN 2008D suggest that the explosion of the core was aspherical \citep{Maund+2009,Gorosabel+2010}. According to a recent theoretical prediction, if such explosion gives rise to a significant deviation from spherical symmetry of the shock at the breakout, it implies a breakout shock velocity that is lower than a spherical breakout velocity, over major parts of the shock front \citep{Matzner+2013}. The good agreement between our spherical model predictions and the observations, and the fact that the breakout velocity we infer, as well as the one estimated by \cite{Soderberg+2008}, are both rather high, as expected for spherical WR shock breakouts, suggest that the asphericity level of the shock at the breakout was low, and its effect on the X-ray emission, as well as the dynamics, was minor. Interestingly, a comparison of the shock velocity and wind density that we find at $R\sim10^{12}$ cm, to that inferred by radio observations at $R\sim3\times10^{15}$ cm, implies that the mass loss rate of the progenitor has increased by more than an order of magnitude during the few days that preceded the SN explosion. A pre-explosion enhanced mass loss may be a common feature of massive stars. \cite{Ofek+2013} reported a mass-loss event just a month before SN 2010mc, likely the explosion of a luminous blue variable progenitor, and suggested a causal connection between the pre-explosion mass loss burst and the explosion. Based on a sample of 16 type IIn SNe, \cite{Ofek+2014} inferred that at least half of the type IIn SNe experience a similar pre-explosion outburst. \cite{Gal-Yam+2014} identified SN 2013cu as an explosion of a WR progenitor and found indications for an increased mass loss rate starting a year before the explosion. Our results for SN 2008D join these accumulating indications and suggest that at least some WR progenitors experience an increased mass loss rate during a short period prior to their explosion. The excess mass loss in this case is due to a higher continuous mass loss, rather than the pre-explosion bursts that may characterize many type IIn SNe. Put together, these findings may indicate on a general phase of enhanced mass loss in the late stellar evolution of massive stars, e.g., by a process as suggested by \cite{Shiode+2014}. G.S. and E.N. were partially supported by an ISF grant (1277/13), an ERC starting grant (GRB-SN 279369), and the I-CORE Program of the Planning and Budgeting Committee and The Israel Science Foundation (1829/12). | 14 | 3 | 1403.3400 |
1403 | 1403.7290_arXiv.txt | We perform a series of simplified numerical experiments to explore how rotation impacts the three-dimensional (3D) hydrodynamics of core-collapse supernovae. For the sake of our systematic study, we employ a light-bulb scheme to trigger explosions and a three-flavor neutrino leakage scheme to treat deleptonization effects and neutrino losses from proto-neutron star interior. Using a $15 \Msun$ progenitor, we compute thirty models in 3D with a wide variety of initial angular momentum and light-bulb neutrino luminosity. We find that the rotation can help onset of neutrino-driven explosions for the models in which the initial angular momentum is matched to that obtained in recent stellar evolutionary calculations ($\sim 0.3$-$3$ rad s$^{-1}$ at the center). For the models with larger initial angular momentum, the shock surface deforms to be more oblate due to larger centrifugal force. This makes not only a gain region more concentrated around the equatorial plane, but also the mass in the gain region bigger. As a result, buoyant bubbles tend to be coherently formed and rise in the equatorial region, which pushes the revived shock ever larger radii until a global explosion is triggered. We find that these are the main reasons that the preferred direction of explosion in 3D rotating models is often perpendicular to the spin axis, which is in sharp contrast to the polar explosions around the axis that was obtained in previous 2D simulations. | Multi-dimensionality in the inner working of core-collapse supernovae (CCSNe) has long been considered as one of the most important ingredients to understand the explosion mechanism. Shortly after the discovery of pulsars in the late 1960's \citep{hewish}, rotation and magnetic fields were first proposed to break spherical symmetry of the supernova engine \citep{leblanc,bisno76,muller79}. In the middle of 1980's, the delayed neutrino-driven mechanism was proposed by \citet{wils85} and \cite{bethe85} who clearly recognized the importance of multi-dimensional (multi-D) effects \citep{wilson88}\footnote{because they obtained more energetic explosions in the spherically symmetric (1D) models when proto-neutron-star (PNS) convection was phenomenologically taken into account.}. Ever since SN1987A, multi-D hydrodynamic simulations have been carried out extensively in a variety of contexts. They give us a confidence that multi-D hydrodynamic motions associated with PNS convection \citep[e.g.,][]{Keil96,mezzacappa98,bruenn04,dessart06}, neutrino-driven convection \citep[e.g.,][]{herant,burr95,jankamueller96,fryer04a,Murphy08}, and the Standing-Accretion-Shock-Instability \citep[SASI, e.g.,][and see references therein]{Blondin03,Ohnishi06,Scheck08,thierry,Iwakami08,Iwakami09,Kotake09,rodrigo09,fern13} can help the onset of neutrino-driven explosions. In fact, a growing number of neutrino-driven models have been recently reported in the first-principle two-dimensional (2D) simulations (\citealt{Buras06b,Buras06a,Ott08,marek,bruenn,bruenn13,suwa10,suwa13,BMuller12b,BMuller13}, however, \citealt{dolence14}). This success, however, is bringing to light new questions. One of the outstanding problems is that the explosion energies obtained in these 2D models from first principles (though some of them were reported before their explosion energies saturated) are typically smaller by one order of magnitudes to explain the canonical supernova kinetic energy ($\sim 10^{51}$ erg, see Tables in \citet{Kotake13} and \citet{Papish14} for a summary). Most researchers are now seeking for some possible physical ingredients to make these underpowered explosions more energetic \citep[see][for recent reviews]{Janka12,Burrows13,Kotake12_ptep}. One of the major candidates is three-dimensional (3D) effects on the neutrino-driven mechanism. Employing a ``light-bulb'' scheme in 3D simulations, \citet{nordhaus10} were the first to point out that 3D leads to easier explosions than 2D. This was basically supported by some of the follow-up studies \citep{burrows12,dolence12,murphy13}, but not by 3D simulations with a similar setup \citep{hanke11,couch12} and by 3D simulations with spectral neutrino transport \citep{Takiwaki12,Hanke13,takiwaki14}. Another prime candidate is general relativity (GR), which has been agreed to help neutrino-driven explosions in both 2D \citep{BMuller12b,BMuller13} and 3D models (\citet{KurodaT12}, see also \citet{Ott12a,kuroda13}). The impacts of nuclear equations of state (EOS) have been investigated in 2D models \citep{marek,Marek09,suwa13,couch13}, and these studies reached an agreement that softer EOS leads to easier explosions. The neutrino-driven mechanism could be assisted by the magnetohydrodynamic (MHD) mechanism that works most efficiently when the pre-collapse cores have rapid rotation and strong magnetic fields (e.g., \citet{Kotake06} and \citet{moesta14} for collective references therein). It should be mentioned that even in the non-rotating case, the MHD mechanism assists the onset of explosion via the magnetic-field amplifications due to the SASI \citep{endeve,endeve12,martin11}. Other possibilities include nuclear burning behind the propagating shock \citep{bruenn06,nakamura12,yamamoto13}, additional heating reactions in the gain region \citep{sumiyoshi08,arcones,furusawa13}, hadron-quark phase transitions \citep[e.g.,][]{sage09,fischer12}, or energy dissipation by the magnetorotational instability \citep{thomp05,Obergaulinger09,masada12,sawai13,sawai14}. Joining in these efforts to look for some possible ingredients to foster explosions, we investigated the roles of rotation in this study. In 2D simulations of a $15 M_{\odot}$ progenitor with detailed neutrino transport, \citet{marek} were the first to observe the onset of earlier explosions in models that include moderate rotation ($\Omega_0 = 0.5~{\rm rad}~{\rm s}^{-1}$ with $\Omega_0$ being the pre-collapse central angular velocity) than in models without rotation. By performing 2D simulations of an 11.2 $M_{\sun}$ star with more idealized spectral neutrino transport scheme \citep[e.g.,][]{idsa}, \citet{suwa10} showed that stronger explosion is obtained in models that include rapid rotation ($\Omega_0 = 2~{\rm rad}$ s$^{-1}$) than those without rotation. \citet{suwa10} also pointed out that rotation helps explosions, not only because the mass enclosed inside the gain radius becomes larger for rapidly rotating models, but also because north-south symmetric bipolar explosions that are generally associated with rapidly rotating models can expel much more material than that of one-sided unipolar explosions in the non-rotating models. These findings naturally open a simple question. Will these features so far obtained in 2D models also persist in 3D models? In this study, we performed a series of simplified numerical experiments to explore how rotation impacts the 3D hydrodynamics of the CCSN core that produces an explosion by the neutrino mechanism. For the sake of our systematic study, we employed a light-bulb scheme to trigger explosions \citep[e.g.,][]{jankamueller96,Janka01} and a three-species neutrino leakage scheme \citep[e.g.,][]{ross03} to treat deleptonization effects and neutrino losses from the neutron star interior (above an optical depth of about unity). It is well-known from interferometric observations \citep[see][for a review]{vanbelle} that stars more massive than 1.5 $M_{\odot}$ are generally rapid rotators \citep{huang06,huang08}. However, due to numerical difficulties of multi-D stellar evolutionary calculations (see \citet{maeder12} for a review, and \citet{meakin07,meakin2011} for recent developments), it is quite uncertain how such high surface velocities are entirely evolved during stellar evolution till the onset of core-collapse, in which multi-D hydrodynamics including mass-loss, rotational mixing, and magnetic braking is playing an active role in determining the angular momentum transport. In this study, we made pre-collapse models by parametrically adding the initial angular momentum to a widely used a 15 $M_{\odot}$ progenitor \citep{WW95}. We carried out 3D special-relativistic simulations starting from the onset of gravitational collapse, through bounce, trending towards explosions (typically up to about $\sim $1 s postbounce) and compare results of thirty 3D models, in which the input neutrino luminosity and the initial rotation rate are systematically varied. We found that a critical neutrino luminosity to obtain neutrino-driven explosions becomes generally smaller for models with larger initial angular momentum. Our 3D models show a much wider variety of the explosion geometry than in 2D. In the 3D rotating models, we found that the preferred direction of explosion is often {\it perpendicular} to the spin axis, which is in sharp contrast to polar explosions around the spin axis that was commonly obtained in previous 2D simulations. We begin in Section \ref{sec:Numerical Method} with a description of numerical setup and initial models. The main results are shown in Section \ref{sec:results}. We summarize our results and discuss their implications in Section \ref{sec:conclusion}. | \label{sec:conclusion} We performed a series of simplified numerical experiments to explore how rotation impacts the 3D hydrodynamics of neutrino-driven core-collapse supernovae. For the sake of our systematic study, we employed a light-bulb scheme to trigger explosions and a neutrino leakage scheme to treat deleptonization effects and neutrino losses from the PNS interior. Using a $15 \Msun$ progenitor, we computed thirty 3D models with a suite of the initial angular momentum and the light-bulb neutrino luminosity. We find that rotation can help the onset of neutrino-driven explosions for models, in which the initial angular momentum is matched to the one obtained in recent stellar evolutionary calculations. For models with larger initial angular momentum, the PNS and the shock surface deform to be more oblate due to the larger centrifugal forces. This makes not only the gain region much more concentrated around the equatorial plane, but also the mass in the gain region bigger. As a result, hot bubbles tend to be coherently formed in the equatorial region, which pushes the shock ever larger radii until the global explosion is triggered. We found that these are the main reasons that the preferred direction of explosion in the 3D rotating models is often perpendicular to the spin axis, which is in sharp contrast to the polar explosions around the axis that was obtained in previous 2D simulations. Finally, as a guide to discuss the strength of our parameterized explosion models, we estimate a {\it diagnostic} energy\footnote{As in \citet{suwa10}, diagnostic energy is defined as the integral of the energy all over zones that have a positive sum of the specific internal, gravitational, and kinetic energy.} and the mass of the PNS (Figure \ref{fig16}). It can be seen that our rapidly rotating models (dotted lines with $\Omega_0 = 0.5\,\pi$, left panel) trending towards an energetic explosion ($\sim 10^{51}$ erg) would possibly leave behind a relatively reasonable remnant mass ($\sim 1.3$-$1.4 \Msun$) for some particular choice of the input luminosity (e.g., dotted green line in the right panel). One has to be careful in interpreting this, because the core neutrino luminosity was assumed to be constant with time and spatially isotropic in this study. It has been shown that neutrino emission from (rapidly) rotating core is not isotropic at all \citep{Janka89_226,Kotake03,Ott08,brandt10}, and more importantly, the neutrino luminosity becomes smaller for models with larger initial angular momentum \citep{Ott08}. Having admitted that the predictive power of our 3D parameterized models especially with rapid rotation is very limited, our 3D models demonstrate that even a {\it moderate} rotation ($\Omega_0 = 0.1\,\pi$) as previously thought, could significantly effect the evolution of the diagnostic energy and the remnant mass (compare solid with dashed line in Figure \ref{fig16}). A self-consistent 3D model \citep[ultimately with 6D Boltzmann transport, e.g.,][]{sumiyoshi12,sumiyoshi14} is apparently needed to have the final word whether rotation will or will not lead to easier onset of neutrino-driven explosions, which we are going to study as a sequel of this work (Takiwaki et al. in preparation). \begin{figure*}[htbp] \plottwo{fig13a.eps}{fig13b.eps} \caption{Time evolution of diagnostic energy (left panel) and the mass of the PNS (defined by a fiducial density of $10^{11}~{\rm g}\,{\rm cm}^{-3}$) for the $L_{\nu,52} = 2.3$ (green lines), $L_{\nu,52} = 2.5$ (red), and $2.7$ (blue) models with different initial rotation rates. For these three values of $L_{\nu,52}$, non-rotating models (dashed lines) do not explode within our simulation time, while moderately rotating (solid) and rapidly-rotating (dotted) models are trending towards explosion. As a reference, horizontal line represents $10^{51}$ erg in the left panel. } \label{fig16} \end{figure*} | 14 | 3 | 1403.7290 |
1403 | 1403.6775_arXiv.txt | { Rotation-powered millisecond radio pulsars have been spun up to their present spin period by a $10^8-10^9{\rm yr}$ long X-ray-bright phase of accretion of matter and angular momentum in a low-to-intermediate mass binary system. Recently, the discovery of transitional pulsars that alternate cyclically between accretion and rotation-powered states on time scales of a few years or shorter, has demonstrated this evolutionary scenario. Here, we present a thorough statistical analysis of the spin distributions of the various classes of millisecond pulsars to assess the evolution of their spin period between the different stages. Accreting sources that showed oscillations exclusively during thermonuclear type I X-ray bursts (nuclear-powered millisecond pulsars) are found to be significantly faster than rotation-powered sources, while accreting sources that possess a magnetosphere and show coherent pulsations (accreting millisecond pulsars) are not. On the other hand, if accreting millisecond pulsars and eclipsing rotation-powered millisecond pulsars form a common class of transitional pulsars, these are shown to have a spin distribution intermediate between the faster nuclear-powered millisecond pulsars and the slower non-eclipsing rotation-powered millisecond pulsars. We interpret these findings in terms of a spin-down due to the decreasing mass-accretion rate during the latest stages of the accretion phase, and in terms of the different orbital evolutionary channels mapped by the various classes of pulsars. We summarize possible instrumental selection effects, showing that even if an unbiased sample of pulsars is still lacking, their influence on the results of the presented analysis is reduced by recent improvements in instrumentation and searching techniques. } | \label{sec:intro} The accretion of $0.1-0.2$ M$_{\odot}$ transferred by a companion star through an accretion disc in a low-mass X-ray binary (LMXB) is the widely accepted process to explain old neutron stars (NSs) that spin at a period of a few milliseconds \citep{alpar1982,radhakrishnan1982,bhattacharya1991}. This framework is known as the {\it recycling} scenario and identifies accreting X-ray bright NSs as the progenitors of radio-millisecond pulsars whose emission is powered by their loss of rotational energy. In this scenario, after a $0.1-10$~Gyr long X-ray-bright accretion phase that spins up the NS \citep[see, e.g.,][]{tauris2012b}, the pressure exerted by the NS magnetic field is able to sweep out the light-cylinder volume, sparking the rotation-powered pulsed radio emission of the {\it recycled} pulsar. Mass-accretion not only spun up the $\ga 300$ binary and isolated rotation-powered millisecond pulsars known in the Galaxy, but probably also reduced their magnetic field strength to values of $10^7-10^9$ G \citep[see, e.g.,][for a review]{bhattacharya1995}. The accretion-driven spin up of an NS in an LMXB is demonstrated by the observation of coherent pulsations at a period of a few ms from 15 NSs (accreting millisecond pulsars, hereafter accreting MSPs; \citealt{wijnands1998}, \citealt{patruno2012c}), and of quasi-coherent-oscillations observed from an additional ten sources exclusively during thermonuclear type I X-ray bursts (nuclear-powered millisecond pulsars, hereafter nuclear MSPs\footnote{We stress that here, the term {\it nuclear} only refers to the mechanism that powers the emission of the source while it shows pulsations, whereas for most of the time, the X-ray output is due to mass accretion.}; \citealt{chakrabarty2003}, \citealt{watts2012}). The swings between rotation and accretion-powered behaviour recently observed from IGR J18245--2452 \citep{papitto2013nat,pallanca2013,ferrigno2013,linares2014} PSR J1023+0038 \citep{archibald2009,stappers2013,patruno2014}, and XSS J12270--4859 \citep{bassa2014,papitto2014,roy2014,bogdanov2014} demonstrated that NSs in some low-mass X-ray binaries alternate cyclically between rotation and accretion-powered states on short time-scales of a few years, or shorter. In this paper, we compare the spin frequency distributions of the various samples of MSPs discovered so far, to test theories for describing the spin evolution of NSs as they evolve from the accretion to the rotation-powered stages of the recycling scenario \citep[see, e.g.,][and references therein]{tauris2012}. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{Freqdistr.eps}} \caption{Spin frequency distribution of binary non-eclipsing and eclipsing rotation-powered millisecond pulsars (top panel; blue and light blue), and nuclear-powered and accreting millisecond pulsars (bottom panel; green and light-green).} \label{fig:Freqdistr} \end{figure} | \label{sec:disc} Our analysis shows that nuclear MSPs form the only class of accretion-powered pulsars that are significantly faster than RMSPs, while accreting MSPs are not. However, if accreting MSPs and eclipsing RMSPs are assumed to form a common class of transitional pulsars, this class is characterized by an spin frequency distribution intermediate between faster nuclear MSPs and slower non-eclipsing radio pulsars. The frequency difference between nuclear MSPs and RMSPs ($\ga 230$ Hz) cannot be due to the weak magneto-dipole spin-down torque that acts onto the NS during the rotation-powered pulsar phase because this develops on a long time-scale, $\tau_{em}\approx 10$~Gyr (for an NS spinning at 500 Hz, with a magnetic dipole of $10^{26}$ G cm$^{3}$). The spin-down torque that developes during the latest stages of the accretion phase, caused by the interaction between the NS magnetic field and the plasma rotating in the disc, is much stronger \citep[up to a factor of $\sim 100$;][]{ghosh1979}. As this torque is more intense when the NS is faster ($\propto \nu^{2}$, e.g., \citealt{rappaport2004}), it eventually balances the rate at which the accreted matter delivers its specific angular momentum to the NS ($\propto \nu^{-1/3}$). This allows us to define an equilibrium period at which the overall torque acting onto the NS vanishes, the exact value of which depends on assumptions on the disc structure and how the field is twisted by the differential rotation of the disc plasma (\citealt{davidson1973,ghosh1979,wang1987,bozzo2009}; see, e.g., \citealt{ghosh2007} for a review). Here, we considered the estimate given by \citet{kluzniak2007} for a 1.4 $M_{\odot}$ NS accreting from an optically thick, gas-pressure-dominated disc, as \begin{equation} \label{eq:magequil} \nu_{\rm eq}^{\rm (magn)}\simeq 723 \: \mu_{26}^{-6/7} \:\dot{M}_{-10}^{3/7}\:\rm{Hz}, \end{equation} where $\mu_{26}$ is the NS magnetic dipole in units of $10^{26}$ G cm$^3$, and $\dot{M}_{-10}$ is the mass accretion rate in units of $10^{-10}$ M$_{\odot}$ yr$^{-1}$. We note that this estimate does not take into account the possible ejection of matter by the rotating magnetosphere (see below). At the end of the accretion stage, the companion star decouples from its Roche lobe over a typical time-scale of $\approx 100$ Myr \citep{tauris2012}. If the NS would spin at the magnetic equilibrium until the complete turn-off of mass accretion, this would result in a spin-down of the NS well below the observed spin frequencies \citep{ruderman1989}. A departure from spin equilibrium must therefore occur at some point. \citet{tauris2012} showed that the spin equilibrium is broken by the rapidly expanding magnetosphere resulting from the significant decrease in the ram pressure of the transferred material when the donor star terminates its mass-transfer phase. The resulting spin-down is aided by the onset of a propeller phase \citep{illarionov1975} that acts to bring the NS farther out of equilibrium during the Roche-lobe decoupling phase. Depending mainly on the magnetic field strength of the accreting pulsar, the loss of rotational energy computed by \citet{tauris2012} typically amounted to about 50~per~cent, corresponding to a decrease in spin frequency of a factor $\sqrt{2}$, when the recycled pulsars switch from the accretor to the final long-term rotation-powered state. This value is of the order of the difference observed between the spin distribution of nuclear MSPs and RMSPs, making the spin-down experienced by the NS during the permanent decoupling of the companion star from its Roche lobe a viable explanation. The spin distribution of accreting MSPs, on the other hand, is closer to that of RMSPs than to nuclear MSPs. This might originate from the differences between the two classes of accretion-powered sources. Unlike nuclear MSPs, accreting MSPs possess a magnetosphere that during the quiescent states may push away the in-flowing matter, letting the NS turn on as a rotation-powered radio pulsar \citep[{\it the radio-ejection mechanism},][]{kluzniak1988,stella1994,campana1998,burderi2001,papitto2013nat}. Because these accreting MSPs spend most of the time in quiescence, their spin evolution is mainly driven by the magneto-dipole torque, and not by stronger torques associated with disc-field interactions. This implies a slower spin evolution than nuclear MSPs, so that a smaller spin difference might be expected between accreting MSPs and their rotation-powered descendants. The only accreting MSP from which many cycles between outbursts and quiescence have been observed, SAX J1808.4-3658 \citep[see][and references therein]{patruno2012b}, spins down at a rate corresponding to a slow-down by just $\sim\!5 $ Hz over the typical time-scale (100~Myr) associated with the Roche-lobe decoupling phase of more massive donor stars that result in helium white dwarf companions. Our analysis did not identify a statistically significant difference between the spin distributions of accreting and nuclear MSPs, although the average frequency shown by the latter class is higher by $\sim\!150$~Hz. However, the differences between these two LMXB classes and the much larger sample of RMSPs can be taken as an indication that a difference between accreting and nuclear MSPs may exist that currently cannot be detected because of the few known objects. A similar conclusion is reached from the significant difference found (95~per~cent confidence level) from comparing nuclear MSPs with the whole sample of transitional MSPs and not only accreting MSPs. If nuclear and accreting MSPs are assumed to spin at or close to the equilibrium value given by Eq.~\ref{eq:magequil}, the contrast between the average frequencies of the two samples can be explained by a difference in the mass-accretion rate of $\Delta\dot{M}/\dot{M}\sim 0.8$, and/or in the dipole moment of $\Delta\mu/\mu\sim -0.4$. Nuclear MSPs generally accrete at a higher average rate than accreting MSPs (see green squares and red circles in Fig.~\ref{fig:mdot}), and most probably possess a weaker magnetic field because they do not show coherent pulsations. A weaker field might also reflect a higher mass-accretion rate as well because a plasma in-flowing at a high rate diamagnetically screens the magnetic field of the NS below its surface \citep{bhattacharya1995,cumming2001,cumming2008}. These considerations may be taken as a clue that nuclear MSPs are found in an earlier evolutionary stage that is characterized by a higher average accretion rate (and a consequently lower dipolar magnetic field) and a faster rotation than to accreting MSPs. The latter are instead observed in a phase much closer to the accretion turn-off and to the switch-on of an RMSP, as also demonstrated by the rapid changes of state observed in some of them. \begin{figure} \includegraphics[height=6.5cm]{mdot.eps} \caption{Long-term mass-accretion rate of millisecond pulsars in the accretion phase, as evaluated by \citet{watts2008}. Accreting-MSPs are plotted as red circles, nuclear MSPs as green squares. Magenta triangles mark the accreting MSPs that showed pulsations only intermittently, and that may have a magnetic field close to the threshold for field burial \citep{cumming2008}. Blue dashed lines represent the magnetic spin equilibrium lines defined by Eq.\ref{eq:magequil}, evaluated for $\mu_{26}=1$ and $\mu_{26}=10$. The red shaded region defines the parameter space in which a source enters the radio ejector state for $\mu_{26}=1$ \citep[e.g.,][]{burderi2001}.} \label{fig:mdot} \end{figure} An important role in explaining the different observed spin distributions of some classes might be played by their different evolutionary history. All accreting- and nuclear MSPs for which the orbital period is known belong to systems with $P_{orb}<1$~day. These binaries originate in systems in which the companion star had a low mass (1-$2\;M_{\odot}$) and was still on the main sequence, at the onset of mass transfer. The same holds for eclipsing RMSPs, that is, the so-called black widows and redbacks, \citep[cf.][and references therein]{chen2013}. In contrast, RMSPs originate in a wider sample of binaries, comprising both systems with low-mass \citep{ts99} and intermediate-mass ($2-6\;M_{\odot}$) companions \citep{tvs00}. Furthermore, depending on their final $P_{orb}$ and the composition of the white dwarf companion, in some of these systems the donor star had already evolved off the main sequence at the onset of the mass transfer. Hence, all RMSPs with massive white dwarf companions and RMSPs with helium white dwarf companions in wide orbits ($P_{orb}\gg 1\;{\rm day}$) are only expected to be mildly recycled as a consequence of the relatively short-lived mass-transfer phase in these systems \citep[see Figs.~1 and 9 in][]{tauris2012b}. Given that the observed sample of non-eclipsing RMSPs contains a fraction of such mildly recycled (relatively slow spinning) sources might therefore partly explain why their observed spin distribution is slightly slower than that of eclipsing RMSPs. However, this selection bias probably does not affect the comparison between nuclear and transitional MSPs because all these known sources evolved from low-mass companions in short ($P_{orb}<1$~day) orbital period systems, ensuring a long phase of mass transfer and efficient recycling. Note that an expected short and intense mass-transfer phase in wide-orbit systems can explain the hitherto non-detection of accreting- and nuclear MSPs in such binaries as well. On the other hand, quantifying the impact of instrumental selection biases is more complex and would deserve a population synthesis analysis, which is beyond the scope of this paper. Nevertheless, some conclusions can be drawn on the basis of the observed properties of MSPs in light of the searching strategies actually employed to detect them. Frequency Doppler shifts induced by the orbital motion limit the time interval $T_{\rm best}$ over which a search for a coherent signal can be performed \citep[$T_{\rm best}\propto P_{orb}^{\,2/3}$; see Eq. 21 in][]{johnston1991}. This significantly reduces the minimum detectable amplitude $A_{min}$ of accretion powered X-ray pulsations emitted by NS in very-short period binaries ($P_{orb}\la$ 1~hr)\footnote{For a 600 Hz pulsar in a $P_{orb}=0.5$~hr binary observed at an X-ray flux of 200 counts s$^{-1}$ (typical, e.g., of observations performed with EPICpn on-board {\textsl XMM-Newton}), $T_{\rm best}$ is reduced below 30 s, a time interval that only allows the detection of strong signals with an amplitude larger than $A_{min}\simeq12$ per cent \citep{vaughan1994}, higher than the typical few per cent amplitude signals observed from accreting MSPs \citep{patruno2012}. A similar result is obtained ($A_{min}\simeq10$ per cent) if a 300 Hz pulsar is considered.}, but since the dependence of this effect on the pulsar spin frequency is only weak ($A_{min}\propto \nu^{1/4}$), the resulting bias on the observed spin frequency distribution of accreting MSPs is small. On the other hand, the orbital motion does not affect the detection of nuclear MSPs at all because the burst oscillations are observed at X-ray fluxes $\ga 100$ times higher than accreting MSPs signals, and a detection is achieved in time-intervals of a few seconds \citep{watts2012}, which is much shorter than the orbital period of the binary. Pulse arrival-time delays induced by orbital motion and by propagation in the ionized interstellar medium also reduce the sensitivity to high-frequency signals in the radio band \citep{dewey1985,hessels2007}. However, recent improvements in data acquisition systems, computational power, and search algorithms have reduced these effects, resulting in doubling the detections of RMSPs in the past five years \citep{lorimer2013,ransom2013}. This increase is also caused by radio follow-up of unidentified $\gamma$-ray {\it Fermi}-LAT sources, which now amount to $15$ per cent of the total number of RMSPs \citep{ray2012,ransom2013}. As the $\gamma$-ray output of a millisecond pulsar increases with its spin-down power ($L_{\gamma}\propto \dot{E}_{sd}^{a}$, with $a=0.5$--$1$ and $\dot{E}_{sd}\propto \mu^2\nu^4$; see \citealt{abdo2013}), the detection of fast-spinning sources through this method is favoured. During the past five years, the follow-up of {\it Fermi} sources determined an eight-fold increase in the number of known Galactic eclipsing RMSP \citep{roberts2013}, which are significantly faster than the rest of RMSP, as we have shown in this work. Before this overall reduction of the bias against the detection of fast radio pulsars, \citet{lorimer2013} estimated the peak of the spin frequency distribution of RMSP of the Galactic plane to be $\Delta\nu\simeq65$ Hz higher than the observed average. This rapid increase of the MSPs population over the past years then indicates that even if a complete sample of fast pulsars has not been achieved yet, the impact of selection bias has been reduced below this level, $\Delta\nu$. In the future, X-ray observatories such as the {\it Large Observatory for X-ray Timing} \citep[LOFT][]{feroci2012}, and to some extent {\it AstroSAT} \citep{agrawal2006} and {\it NICER} \citep{gendreau2012}, are expected to increase the sensitivity to fast coherent signals from accreting sources up to a factor $\sim 5$. At the same time, further improvement in the computational power and analysis techniques used to discover RMSPs, the numerous candidates identified by {\it Fermi}-LAT, and the advent of the {\it Square Kilometre Array} \citep[SKA][]{smits2009} will also increase the number of fast radio pulsars in close binary systems. An increase in the number of sources known will be fundamental to constrain the details of pulsar recycling from the point of view of the observed spin distributions. | 14 | 3 | 1403.6775 |
1403 | 1403.6891_arXiv.txt | The relative abundances of carbon and oxygen have long been recognized as fundamental diagnostics of stellar chemical evolution. Now, the growing number of exoplanet observations enable estimation of these elements in exoplanetary atmospheres. In hot Jupiters, the C/O ratio affects the partitioning of carbon in the major observable molecules, making these elements diagnostic of temperature structure and composition. Here we present measurements of carbon and oxygen abundances in 16 stars that host transiting hot Jupiter exoplanets, and compare our C/O ratios to those measured in larger samples of host stars, as well as those estimated for the corresponding exoplanet atmospheres. With standard stellar abundance analysis we derive stellar parameters as well as [C/H] and [O/H] from multiple abundance indicators, including synthesis fitting of the [O I] 6300\,{\AA} line and NLTE corrections for the O I triplet. Our results, in agreement with recent suggestions, indicate that previously-measured exoplanet host star C/O ratios may have been overestimated. The mean transiting exoplanet host star C/O ratio from this sample is 0.54 (C/O$_{\odot}$=0.54), versus previously-measured C/O$_{\rm{host~star}}$ means of $\sim$0.65-0.75. We also observe the increase in C/O with [Fe/H] expected for all stars based on Galactic chemical evolution; a linear fit to our results falls slightly below that of other exoplanet host stars studies but has a similar slope. Though the C/O ratios of even the most-observed exoplanets are still uncertain, the more precise abundance analysis possible right now for their host stars can help constrain these planets' formation environments and current compositions. | To date, the most statistically significant trend in host star abundances pertains to metallicity. Stars hosting giant, close-in planets have higher metallicities (measured as [Fe/H]\footnote{[X/H]=log$N$(X) - log$N$(X)$_{\odot}$, where log$N$(X)=log$N$(X/H)+12}) than stars without detected giant planets (e.g., Gonzalez 1998 \& 2001; Santos et al.\,2004; Fischer \& Valenti 2005; Ghezzi et al.\,2010). Statistical studies of dwarf stars hosting planets indicate a metallicity enhancement of $\sim$0.15 dex for stars with giant planets and a 99.9994\% probability that stars with/without giant planets are drawn from different parent populations (Buchhave et al. 2012; Ghezzi et al. 2010). However, the host star metallicity trend is weaker for Neptune-sized planets -- the difference in the mean [Fe/H] of Jovian-mass hosts versus Neptunian-mass hosts is $\sim$0.10 dex, with the Neptune-mass hosts showing lower [Fe/H] values (e.g., Ghezzi et al.\,2010). Smaller planet (R$_{\rm{P}}\leq$4 R$_{\oplus}$) host stars from the \textit{Kepler} sample show no metallicity enhancement, and have a flatter distribution of metallicities, though roughly peaked at solar. These smaller planet host stars have a probability between 0.98 and 0.9996 of originating from a different parent population as larger planet host stars from \textit{Kepler} (Buchhave et al.\,2012; Everett et al.\,2013). Looking beyond the correlation between planet size and stellar metallicity, several studies have searched for other trends between planet parameters and host star abundances indicative of planet formation conditions. Mel\'endez et al.\,(2009) find through abundance analyses of 11 solar ``twins'' that the Sun is deficient by $\sim$20\% in refractory elements, which have condensation temperatures $T_c\gtrsim$ 900 K, relative to volatile elements when compared to other solar twins. This trend of decreasing refractory elemental abundances as a function of $T_c$ is suggested to be a signature of terrestrial planet formation -- the ``missing'' refractory elements from the stellar photosphere are incorporated into rocky planets (Mel\'endez et al.\,2009). However, subsequent studies of similar precision measurements of solar analogs (Gonz\'alez Hern\'andez et al.\,2010; Gonz\'alez Hern\'andez et al.\,2013) and stars with planets (Schuler et al.\,2011a \& 2011b) across a range of $T_c$ show that the abundance patterns of stars with and without planets are not significantly different, or may be indistinguishable from Galactic chemical evolution effects. New evidence from Jupiter- and Neptune-sized planet host stars that are more metal-rich, or warmer than the Sun and have less massive convective envelopes, indicates that the depletion signature may depend on the stellar convective envelope size at the time of planet formation, and thus the timescale of disk dispersal around different types of stars (Ram\'{\i}rez et al.\,2013) The growing number of transiting and directly-imaged exoplanet observations enable estimates of elemental and molecular abundances in the atmospheres of the planets themselves. The \textit{Hubble Space Telescope} and \textit{Spitzer Space Telescope}, aided by multiple ground-based facilities, have detected the most abundant molecules (H$_2$O, CO, CH$_4$, CO$_2$) in the atmospheres of several of the brightest transiting planets (e.g., Tinetti et al.\, 2007; Swain et al.\,2008, 2009ab; Snellen et al.\,2010; Beaulieu et al.\,2010). Ground-based observatories have made similar strides in studying the molecular properties of a handful of directly imaged self-luminous exoplanets (e.g., Marois et al.\,2008; Barman et al.\,2011ab; Skemer et al.\,2012, 2013; Konopacky et al.\,2013; Janson et al.\,2013). The differences and trends between the elemental compositions of host star and exoplanet atmospheres provide clues about the formation and evolution processes of planetary systems. \subsection{The Role of Carbon and Oxygen} Carbon and oxygen are important players in the composition of stars and planets, as the third and fourth most abundant elements in the universe. The measurement of C and O in stars, especially with respect to iron, which is produced in both Type Ia and Type II supernovae, serves as a fundamental diagnostic of the chemical enrichment history of the Galaxy. The impact of massive stars' Type II supernovae, and thus the major oxygen contributor, lessens with time and increasing metallicity as the influence of low- and intermediate-mass stars' carbon contribution grows. Measuring C and O in exoplanets is diagnostic of current atmospheric composition and temperature structure: the atmospheric C/O ratio\footnote{The C/O ratio -- the ratio of carbon atom to oxygen atoms -- is calculated in stellar abundance analysis as C/O= N$_{\rm{C}}$/N$_{\rm{O}}$=10$^{\rm{logN(C)}}/10^{\rm{logN(O)}}$.} affects the molecular composition, and hence observed spectral signatures, through thermochemial equilibrium partitioning of carbon in CO, CH$_{4}$, and CO$_2$. The C/O ratio can also reflect where in the protoplanetary disk a planet formed, as well as subsequent migration and evolution (e.g., Stevenson \& Lunine 1988; Gaidos 2000; Ciesla \& Cuzzi 2006; \"Oberg et al.\,2011). Theories of planet formation describe how close-in giant planets form in the outer protoplanetary disk , where icy planetesimals coalesce into a core, which accretes gas and migrates inwards (e.g., Pollack et al.\,1996; Owen et al.\,1999; Ida \& Lin 2004b). The main molecular reservoirs of C and O have different condensation temperatures ($T_c$), so their relative amounts vary at different temperatures and disk radii, as do the amounts of these molecules in gas or solid form (\"Oberg et al.\,2011). Gas and grains also move differently in the disk with time, as grains grow and decouple from the gas, sequestering solid material beyond the ``ice'' lines of different molecules (e.g., Ciesla \& Cuzzi 2006; \"Oberg et al.\,2011). The C/O ratio of a planet therefore does not necessarily reflect the protoplanetary-disk-averaged C/O ratio, and instead may point towards localized concentrations/depletions of carbon- and oxygen-bearing molecules (Ciesla \& Cuzzi 2006; \"Oberg et al.\,2011; Najita et al.\,2013). Many groups have performed stellar abundance analyses of exoplanet host stars in order to determine their physical parameters (T$_{\rm{eff}}$, log \textit{g}, [Fe/H]) and chemical abundances, to further study the trends discussed above (e.g., Delgado Mena et al.\,2010; Petigura \& Marcy 2011; Brugamyer et al.\,2011; Schuler et al.\,2011ab; Nissen 2013). However, only a few transiting exoplanet host stars have published abundances other than [Fe/H] or the more generic [M/H] (e.g., HD 209458, Schuler et al.\,2011a; WASP-12, Petigura \& Marcy 2011; 55 Cnc, Bond et al.\,2010 \& Teske et al.\,2013b; XO-2, Teske et al.\,2013a). Here we add to the small sample of transiting exoplanet host stars with measured abundances beyond [Fe/H], and the handful with measured C/O ratios. We report on sixteen transiting hot Jupiters hosts to investigate the extent to which we can relate host star compositions to those of their planets, and search for carbon-rich planet formation environments. The sample presented here contains the host stars of some of the most-observed exoplanets whose atmospheres can and are being modeled to constrain their C/O ratios (e.g., Madhusudhan 2012; Moses et al.\,2013; Line et al.\,2013). This work provides a step toward comparing specific host star and exoplanet atmospheres to search for the chemical effects of exoplanet formation. | Our final adopted stellar parameters and their 1$\sigma$ uncertainties for each target are listed in Table \ref{tab:stellar_params}, and the adopted elemental abundances and their 1$\sigma$ uncertanties for each target are listed in Table \ref{tab:compare}. We compared our results to those of the catalog of Stars With ExoplanETs (SWEET-Cat), described in Santos et al.\,(2013), and those determined by Torres et al.\,(2012; T12). These two references are comparable to ours in their analysis methods and samples (transiting exoplanet host stars). SWEET-Cat compiles sets of atmospheric parameters previously published in the literature and, whenever possible, derived using the same uniform methodology of Santos et al.\,(2004). The main sources of stellar parameters in SWEET-Cat for the targets in our sample are Santos et al.\,(2004), Ammler-von Eiff et al.\,(2009), and Mortier et al.\,(2013). SWEET-Cat reports T$_{\rm{eff}}$, [Fe/H], and log $g$ for all of the targets in our sample of stars, and $\xi$ for ten of the targets. T12 compared the resulting T$_{\rm{eff}}$ and [Fe/H] from three different stellar analysis programs. They include stellar parameter classification (SPC; Buchhave et al.\,2012), Spectroscopy Made Easy (SME; Valenti \& Piskunov 1996), and MOOG, the latter of which we employ here. They also determine log $g$ values, but do not report them, and $v$ sin $i$ values, which we do not determine in our work. T12 report final ``averaged'' T$_{\rm{eff}}$ and [Fe/H] values from all attempted analysis methods for thirteen of the stars in our sample, though MOOG-derived T$_{\rm{eff}}$ and [Fe/H] values are included in that average for only ten of the stars in our sample. In all cases our derived [Fe/H] values are consistent with those of SWEET-Cat and T12-average, within uncertainties. The median $\Delta$[Fe/H] (as in $|$ours - theirs$|$) values are 0.08 and 0.02 for SWEET-Cat and T12-average, respectively. The T$_{\rm{eff}}$ and log $g$ values reported here also overlap within uncertainties the values reported in T12-average and those in SWEET-Cat in almost every case. The median $\Delta$log $g$ for SWEET-Cat is 0.17, and the median $\Delta$T$_{\rm{eff}}$ for SWEET-Cat and T12-average are 55 K and 41 K, respectively. In the two cases where T$_{\rm{eff}}$ does not overlap SWEET-Cat our values are cooler by 54 K (WASP-12) and 202 K (WASP-32). Two other sources (Brown et al.\,2012; Maxted et al.\,2010) report a very similar T$_{\rm{eff}}$ for WASP-32 as the (cooler) value derived here, and the results of Torres et al.\,(2012) for WASP-12 agree well with our T$_{\rm{eff}}$. In the three cases where log $g$ does not overlap SWEET-Cat or T12-average our values differ % by $\leq$0.04 dex. \subsection{Comparison to Previous Studies of C and O in Exoplanet Host Stars} This study focuses on transiting exoplanet host star elemental abundances, particularly their C/O ratios. No previous study of which we are aware has uniformly derived [C/H], [O/H], and C/O values for these stars. However, several studies have examined C/O ratios in non-transiting exoplanet host stars versus stars not known to host planets. [Any star designated as a ``non-host'' has the potential to harbor a smaller ( undetected) planet; indeed, it may be the case that most stars have one or more small planets (e.g., Cassan et al.\,2012).] Here we compare our results to these other host star C/O studies. Bond et al.\,(2006) measure [C/H] in 136 G-type stars, 20 of which are exoplanet hosts, and Bond et al.\,(2008) measure [O/H] ratios in 118 F- and G-type stars stars, 27 of which are known exoplanet hosts. Line lists are not explicitly given for the measured C lines in Bond et al.\,(2006); Bond et al.\,(2008) use the high-excitation O I triplet at $\lambda$ = 7771.9, 7774.2, and 7775.4\,{\AA}. Bond et al.\,(2010) also compiled C/O ratios derived from measurements in Ecuvillion et al.\,(2004) and (2006). Ecuvillion et al.\,(2004) measures [C/H] from the two lowest excitation lines of carbon (5052.17 and 5380.34\,{\AA}), and Ecuvillion et al.\,(2006) measures [O/H] from the forbidden [O I] line at 6300.3 \AA, the high-excitation O I triplet at $\lambda$ = 7774 \AA, and a set of 5 near-UV OH lines around 3100 \AA. Both of the Ecuvillion studies and the Bond studies implement an analysis method similar to that performed here, with the spectral synthesis code MOOG and a grid of Kurucz (1993) ATLAS9 model atmospheres, although these studies do not derive the host star parameters (T$_{\rm{eff}}$, log \textit{g}, $\xi$, and [Fe/H]), only specific elemental abundance ratios. In the host star sample reported in Bond et al.\,(2010), 35\% have C/O$>$0.8 (they did not specifically report any non-host stars); Bond et al.\,(2008) find 8\% of host stars and 5\% of non-host stars in their sample to have C/O$>$0.8. Following the Bond et al. investigation, two larger studies of the C/O ratios of non-host stars versus host stars were conducted. Delgado Mena et al.\,(2010) measure carbon and oxygen in 100 host stars, along with 270 non-host stars, using the C I lines at 5052.17 {\AA} and 5380.34\,{\AA} and the [O I] forbidden line at 6300.3 {\AA}. They measure equivalent widths with the ARES program\footnote{The ARES code can be downloaded at http://www.astro.up.pt/∼sousasag/ares/.} (Sousa et al.\,2007), and used MOOG and Kurucz ATLAS9 model atmospheres (Kurucz\,1996) for abundance analysis. Delgado Mena et al.\,(2010) find 34\% of their measured host stars have C/O$>$0.8, while in their non-host sample the fraction of stars with C/O$>$0.8 is 20\%. Petigura \& Marcy\,(2011) find carbon and oxygen abundances for 704 and 604 stars, respectively, but only 457 have reliable measurements for both elements that can be used to determine C/O ratios, 99 of which are exoplanet hosts. These authors measure the 6587 {\AA} C I line for carbon and the [O I] line at 6300.3 {\AA} for oxygen, and use the SME code with Kurucz (1992) stellar atmospheres for their abundance analysis. Petigura \& Marcy\,(2011) find % 34\% of host stars in their sample have C/O$>$0.8, versus 27\% of non-host stars in their sample with with C/O$>$0.8. The goal of this paper is to investiage and constrain values of stellar host C/O ratios in systems with observed transiting giant planets, since transit spectroscopy potentially allows for determinations of the corresponding planetary C/O ratios. This goal is driven partially by the recent suggestion by Fortney\,(2012) that the C/O ratios of both host- and non-host stars in the studies noted above have been overestimated due to errors in the derived C/O ratios and the observed apparent frequency of carbon dwarf stars implied by these studies. Nissen\,(2013) recently rederived the carbon and oxygen abundances for 33 of Delgado Mena et al.\,(2010)'s host stars that have additional ESO 2.2m FEROS spectra covering the O I triplet at 7774\,{\AA}, which was not originally used by Delgado Mena. He implements a differential analysis with respect to the Sun, with equivalent widths of C and O measured in IRAF with Gaussian profiles and abundances derived by matching the observed equivalent widths with those measured in plane parallel MARCS atmosphere models (Gustafsson et al.\,2008) having the same stellar parameters as those published by Delgado Mena et al.\,(2010). Accounting for NLTE effects on the triplet line strengths by using the Fabbian et al.\,(2009) corrections, Nissen\,(2013) finds differences from Delgado Mena et al.\,(2010) in the derived oxygen abundances. This results in both a tighter correlation between [Fe/H] and C/O (Nissen finds C/O = 0.56$+$0.54[Fe/H] with an rms dispersion $\sigma$(C/O)=0.06), as well as a much smaller fraction of host stars with C/O$>$0.8 (only 1 out of 33). The tight trend of increasing C/O with [Fe/H], e.g., Nissen (2013; as noted above), is indicative of the importance of overall Galactic chemical evolution in setting the fraction of dwarf stars that might be carbon rich. The increase in C/O with metallicity points to the importance of low- and intermediate-mass star carbon nucleosynthesis at later, more metal-rich times. The influence of mass-star Type II supernovae, the major oxygen contributors, is diluted with time as low- and intermediate-mass stars become more important, thus C/O increases. The fraction of ``carbon-rich'' (C/O $\geq$0.8) planet-hosting stars is thus expected to increase with increasing metallicity in the disk, with the Nissen trend indicating metallicities greater than [Fe/H]$\sim$0.4 might begin to have significant fractions of carbon-rich dwarf stars. All of the planet-hosting stellar samples discussed here have very few, if any, stars at these metallicites or higher. This paper differs from the studies listed above because 1) the sample here is much smaller, being limited to only hosts of transiting exoplanets, and 2) only one non-host star is included (the binary companion XO-2S). In Figure \ref{comparison_plot1} the host star abundances derived here, shown with red filled circles with error bars, for [C/H] and [O/H] versus [Fe/H] are compared to the results of the large samples of Delgado Mena et al.\,(2010), shown as gray asterisks for host stars and open squares for non-host stars, and Nissen\,(2013), shown as blue asterisks. All three studies define similar behaviors of [C/H] and [O/H] as a function of [Fe/H], with the carbon exhibiting larger slopes with iron relative to oxygen; this illustrates the increasing importance of carbon production from low- and intermediate-mass stars relative to massive stars with increasing chemical maturity. The slopes of the trends in Figure \ref{comparison_plot1} are all quite similar, with the Nissen (2013) trend exhibiting the smallest scatter about a linear fit. Using the results of this paper, linear trends are fit to both [C/H] and [O/H] versus [Fe/H] with the following results: [C/H] = 0.95[Fe/H] - 0.05 and [O/H] = 0.56[Fe/H] + 0.01. Excluding the apparently carbon-rich outlier HD 189733, these fits are [C/H] = 1.02[Fe/H] - 0.08 and [O/H] = 0.56[Fe/H] + 0.01; we refer to these fits without HD 189733 throughout the rest of the paper. Quantitatively, the C versus Fe slope is about twice as large (in dex) as that for O versus Fe based on the linear fits to the abundances derived here. The relation for carbon passes 0.08 dex below solar (i.e., at [Fe/H],[C/H]=0,0) while passing close to solar for oxygen, offset by only +0.01 dex. Because of these even rather small offsets ($\sim$0.05-0.1 dex), the C/O ratios as a function of [Fe/H] might fall below solar as defined by our results for this particular sample of stars: an offset of -0.05 to -0.1 dex in [C/O] would correspond to an offset of 0.1 to 0.2 lower in a linear value of C/O. This does not necessarily correspond to simply errors in the analysis, but may reflect both fitting linear relations to our results, which are probably only approximate descriptions of the real Galactic disk relation, as well as there not being a universal trend of [C or O]/H versus [Fe/H]. The offsets here most probably reflect both uncertainties in the analysis (already discussed in Section 3.2) and intrinsic scatter in real Galactic disk populations that will map onto the sample of stars analyzed here. Figure \ref{comparison_plot2} illustrates values of C/O versus [Fe/H] from this study, along with those values from Delgado Mena et al.\,(2010) and Nissen (2013). All three studies find a clear increase in C/O verus [Fe/H], which represents the signature of Galactic chemical evolution, as discussed previously. The relation in C/O in this work falls somewhat below those of the other two studies, but all three exhibit similar slopes. This similarity is born out by a quantitative comparison of C/O versus [Fe/H] between Nissen (2013) and this study. Nissen derived a linear fit of C/O = 0.54[Fe/H] + 0.56, while the same fit to the results here find C/O = 0.53[Fe/H] + 0.45. The adopted solar value in this study of C/O = 0.54 is larger by 0.09 than the value defined by the best-fit linear relation defined by our sample of stars and our results. Including the outlier HD 189733 in our linear fit results in C/O = 0.43[Fe/H] + 0.49, corresponding to a C/O ratio 0.05 smaller than the solar value of 0.54; including this outlier increases the scatter around the fit from 0.04 to 0.1. A linear difference of 0.09 corresponds to 0.08 dex for the solar-relative [C/O], which is comparable to the repsective offsets of -0.08 dex and +0.01 dex in the [C/H] and [O/H] relations. Another way of investigating the inherent scatter within our results is to remove the linear best-fit from the values of C/O and look at the scatter about the fitted relation. When this is done the median residual scatter in C/O is $\pm$0.04 , or 0.03 dex in [C/O]. This comparison of C/O versus [Fe/H] trends between Nissen (2013) and this study indicates that the derived slopes are very similar, but there remain small offsets in zero-point C/O of $\sim$ 0.10 - 0.15 caused by a combination of differences in (presumably) the stellar samples, the adopted solar C/O ratios (0.58 for Nissen and 0.54 for this study), as well as the abundance analysis, e.g., much of the offset is due to somewhat smaller values of [C/H] at our lower metallicity range. The mean and standard deviation of the [C/H], [O/H], and C/O distributions from this study, as well as those from all the previous studies of host star carbon and oxygen abundances mentioned above, are listed in Table \ref{tab:compare2}. The mean [C/H] of the transiting exoplanet hosts in this paper is less than the mean [C/H]$_{\rm{host}}$ from the previous works, 0.14 in the five previous studies versus 0.08 found here. The mean [O/H] value found in our sample is the same as the mean [O/H]$_{\rm{host}}$ from previous studies, 0.07. However, the standard deviations of [C/H]$_{\rm{transiting}}$ and [O/H]$_{\rm{transiting}}$ from this paper are large, 0.20 and 0.13, respectively, so any differences in our mean [C/H] and [O/H] values are to be viewed with caution. The mean C/O ratio of the transiting exoplant host stars in our sample is 0.54, with a standard deviation of 0.15, versus the mean from the previous papers of 0.71 with a standard deviation of 0.07. Therefore, the sample of carbon and oxygen abundance ratios for transiting exoplanet host stars presented here, while marginally consistent, are on average lower than those measured by other groups for non-transiting exoplanet host stars. Our measurements are more in line with the suggestions by Fortney\,(2012) and Nissen (2013) that prior studies overestimated C/O ratios; the mean C/O$_{\rm{hosts}}$ of Nissen is 0.63$\pm$0.12. However, as noted by Fortney\,(2012), each previous study scales their C/O ratios based on different log$N$(C)$_{\odot}$ and log$N$(O)$_{\odot}$ values. Delgado Mena et al.\,(2010) list log$N$(C)$_{\odot}$ and log$N$(O)$_{\odot}$ as 8.56 and 8.74, respectively, resulting in C/O$_{\odot}$=0.66. These are also the values listed in Ecuvillion et al.\,(2004) and (2006), the quoted sources of Bond et al.\,(2010). Petigura \& Marcy\,(2011) list log$N$(C)$_{\odot}$ and log$N$(O)$_{\odot}$ as 8.50 and 8.70, respectively, resulting in C/O$_{\odot}$=0.63. Nissen\,(2013)'s C/O$_{\odot}$=10$^{8.43}$/10$^{8.665}$=0.58. Figure \ref{comparison_plot2} illustrates the different C/O$_{\odot}$ from Delgado Mena et al.\,(2010) and Nissen\,(2013). Accounting for the difference in log$N$(C)$_{\odot}$ and log$N$(O)$_{\odot}$ decreases the average C/O ratios from the other sources (from top to bottom) in Table \ref{tab:compare2} by $\sim$0.15, $\sim$0.13, $\sim$0.11, and $\sim$0.05, closer to the average C/O ratio we derive for our sample. Figure \ref{comparison_plot2}'s right panel shows the C/O ratios of Delgado Mena et al.\,(2010) and Nissen\,(2013) along with those derived in this work, all on the same scale, illustrating how using different solar C and O absolute abundances changes the resulting C/O ratios. This underscores the caution, as mentioned in Fortney\,(2012) and Nissen\,(2013), required when directly comparing C/O ratios derived from different groups. We now focus on the C/O ratios in each studied system to investigate possible links between host star C/O ratios with planetary and system properties. \subsection{Trends with C/O$_{\rm{host\,star}}$ versus Planetary Parameters} Presently there are two major observed trends relating stellar chemical composition to the presence of planets -- hot Jupiter exoplanets are more often found around intrinsically higher-metallicity stars (e.g., Fischer \& Valenti 2005), and the fraction of stars with giant planets increases with stellar mass (e.g., Johnson et al.\,2010; Ghezzi et al.\,2010; Gaidos et al.\,2013). Measuring potential host stars' chemical abundance distributions may develop into a powerful tool for inferring the presence, or even specific type (size, orbit, composition), of exoplanets around different types of stars. This technique is of increasing importance in the context of large surveys that are discovering exoplanets, and targeted studies of unusual or potentially-habitable exoplanets. In this study we explore whether the stellar C/O ratio has predictive power with respect to hot Jupiter properties, particularly the exoplanetary atmosphere compositions. Characteristic observations of the atmospheres of the exoplanets in this sample -- the \textit{Spitzer/IRAC} 3.6, 4.5, 5., and 8.0 $\mu$m secondary eclipse fluxes -- as well as their physical properties like mass, radius, semi-major orbital axis, period, and equlibrium temperature were gathered from the NASA Exoplanet Archive\footnote{http://exoplanetarchive.ipac.caltech.edu/} and compared to host star C/O ratios. By eye it appears that planet radius and planet equlibrium temperature may decrease with increasing C/O$_{\rm{host\,star}}$ (Fig.\ref{pparam}), but these trends are dominated by one or two points and, once these points are removed, no significant trends with planet parameters are found. We also find weak negative correlations between each system's C/O$_{\rm{host star}}$ and planetary \textit{Spitzer}/IRAC secondary eclipse fluxes (e.g., $r\sim$-0.4 to -0.6), but these correlations are not statistically significant ($p>$0.05). This lack of trends between C/O$_{\rm{host\,star}}$ is perhaps not surprising. The hot Jupiter host stars in this sample were chosen based on the amount of observational data that exists for their planets, and thus how ``characterizable'' their planets' atmospheres are, with the goal of directly comparing star and planet C/O ratios. No planetary or stellar parameters serve as ``control variables'' in this study, and our sample is actually diverse in both respects. The host stars span 5100$\lesssim$T$_{\rm{eff}}\lesssim$6470 K, -0.21$\lesssim$[Fe/H]$\lesssim$0.44, and spectral types F6 through K1. The planets in our sample range in mass from $\sim$0.5-4.2 M$_{\rm{J}}$, in period from $\sim$1.09-111 days (with the second longest being 4.5 days), in density from $\sim$0.2-8 g cm$^{-3}$ (with the second most-dense being 3.4 g cm$^{-3}$), and in equilibrium temperature from $\sim$400-2500 K. That we do not find a significant correlation between C/O$_{\rm{host\,star}}$ and any of these planetary parameters implies that (1) our sample may yet be too small to reveal distinct trends, and/or (2) the influence of the host star C/O ratio is a more complex function of multiple parameters of the planet and/or its formation history. While (1) is possible, (2) also seems likely and could result in the C/O ratio comparison between stars and planets serving a more interesting function. In protoplanetary disks, different condensation fronts due to temperature and the movement of gas and grains in the disk can change the relative ratios and/or of carbon and oxygen as compared to those in the parent star (Stevenson \& Lunine 1988; Lodders 2010; Ciesla \& Cuzzi 2006; \"Oberg et al. 2011). In particular, the enhancement or depletion of water and thus oxygen is sensitive to the size and migration of icy solids in the disk, so the C/O ratios of the inner and outer disk regions evolve with time and depend on both initial conditions and the efficiency with which solids grow to large sizes (Ciesla \& Cuzzi 2006; Najita et al.\,2013). Overall, the final C/O$_{\rm{planet}}$ does not necessarily reflect the C/O$_{\rm{disk-average}}$, and depends % on the location and timescale of formation, how much of the atmosphere is accreted from gas versus solids, and how isolated the atmosphere is from mixing with core materials (Ciesla \& Cuzzi 2006; \"Oberg et al.\,2011). In our own solar system gas giant planets, oxygen is not well constrained because water, the major oxygen carrier, condenses deeper down in their cool (T$\leq$125 K) atmospheres, out of the observable range of remote spectra (Madhusudhan 2012). However, carbon is known to be enhanced above solar by factors of $\sim$2-6, 6-11, 18-50, and 28-63 in Jupiter, Saturn, Uranus, and Neptune, respectively (Wong et al.\,2008 and references therein). Thus, though the composition of the host star provides a good estimate of the system C/O ratio and the natal molecular cloud environment, differences between the host star and planetary C/O ratios may be common, and may be used to probe where and when in the disk the planet formed. A third possibility is that the host star C/O ratio has no connection to the formation of planets and is not a useful metric for distinguishing planet types. However, theoretical results (e.g., Johnson et al.\,2012; Ali-Dib et al.\,2013) demonstrating the influence of the host star C/O ratio on the composition of protoplanetary disk, and recent observations (e.g., Najita et al.\,2013; Favre et al.\,2013) indicating that disks themselves likely have a range of C/O ratios which are related to other planet formation parameters (mass of the disk, grain growth and composition, etc.), suggest that C/O ratios of host stars do play a role, at some stage, in planet formation. \subsection{Carbon and Oxygen in Specific Exoplanet Systems} A small fraction of exoplanets, mostly hot Jupiters orbiting very close to their host stars, have been observed and analyzed with spectroscopy and photometry in the optical and near-infrared during primary transit (e.g., Madhusudhan \& Seager 2009; Swain et al.\,2008, 2013; Moses et al.\,2011; Mandell et al.\,2013) and/or secondary eclipse (e.g., Charbonneau et al.\,2005, 2008; Knutson et al. 2008; Madhusudhan et al.\,2011; Crossfield et al.\,2012). Direct imaging of exoplanets in wider orbits (e.g., Marois et al.\,2008 \& 2010; Lagrange et al.\,2009; Bailey et al.\,2014) has also opened up for study a new population of self-luminous planets in Jovian-type orbits. As discussed in the introduction, a gas giant planet's C/O ratio has important implications for its composition. At the temperatures and pressures characteristic of such atmospheres, a high C/O ratio ($\gtrsim$0.8) can significantly alter the temperature and chemistry structure by depleting the dominant opacity source H$_2$O and introducing new sources that are C-rich like CH$_4$, HCN, and/or other hydrocarbons. In thermochemical equilibrium, C/O$>$1 causes O to be confined mostly to CO, depleting H$_2$O and enhancing CH$_4$ versus what is expected in solar-abundance atmospheres (C/O$_{\odot}$=0.55$\pm$0.10; Asplund et al.\,2009; Caffau et al.\,2011), which have abundant H$_2$O and CO (Madhusudhan 2012; Moses et al.\,2013). In carbon-rich atmospheres, the temperature controls how depleted the H$_2$O is compared to solar and the partitioning of carbon between CH$_4$ and CO, which in turn influences the oxygen balance between CO and H$_2$O (Madhusudhan 2012). With the ability to constrain exoplanet atmosphere compositions (e.g., Madhusudhan 2012; Lee et al.\,2012; Moses et al.\,2013; Konopacky et al.\,2013; Line et al.\,2013), a logical next step towards determining the host star's influence on exoplanet formation is the direct comparison of the abundance ratios of star/planet pairs. \subsubsection{WASP-12} For WASP-12b, one of the brightest transiting exoplanets, the comparison between host star and planet composition has already begun (Madhusudhan et al.\,2011; Madhusudhan 2012; Petigura \& Marcy 2011; Crossfield et al.\,2012; Swain et al.\,2013; Copperwheat et al.\,2013; Sing et al.\,2013). The host star is found in this work to have [Fe/H] = 0.06$\pm$0.08 and C/O=0.48$\pm$0.08. We note that this metallicity differs significantly from the [M/H]= 0.30$^{+0.05}_{-0.10}$ reported by Hebb et al.\,(2009) in the WASP-12b discovery paper, based on spectral synthesis of four regions including the Mg b triplet at 5160-5190\,{\AA}, Na I D doublet at 5850-5950\,{\AA}, 6000-6210 \AA, and H$\alpha$ at 6520-6600\,{\AA}, following the procedure of Valenti \& Fischer\,(2005). Coupling their atmospheric modeling and retreival methods to published secondary eclipse photometry and spectroscopy spanning 0.9 $\mu$m to 8 $\mu$m, Madhusudhan et al.\,(2011) and Madhusudhan\,(2012) suggest that WASP-12b's atmosphere has a C/O ratio$\geq$1. Their best fit describes an atmosphere abundant in CO, depleted in H$_2$O, and enhanced in CH$_4$, each by greater than two orders of magnitude compared to the authors' solar-abundance, chemical-equilibrium models. However, this high C/O$_{\rm{planet}}$ ratio for WASP-12b's atmosphere is ruled out at the $>$3$\sigma$ level with new observations at 2.315 $\mu$m and reanalysis of previous observations accounting for the recently detected close M-dwarf stellar companion (Bergfors et al.\,2011; Crossfield et al.\,2012). Including the dilution of the reported transit and eclipse depths due to the M-dwarf, the dayside spectrum of WASP-12b is best explained by a featureless 3000 K blackbody (Crossfield et al.\,2012). Subsequent data (Sing et al.\,2013) do not detect metal hydrides MgH, CrH, and TiH or any Ti-bearing molecules, which were previously suggested as indicative of high-C/O ratio scenarios (Madhusudhan 2012; Swain et al.\,2013). A C/O$<$1 composition for WASP-12b is also consistent with the study of Line et al.\,(2013), who use a systematic temperature and abundance retrieval analysis, combining differential evolution MCMC with an optimal-estimation-based prior, to rule out strong temperature inversion in WASP-12b's atmosphere and thus the presence of TiO causing such an inversion. Accounting for the M dwarf companion, these authors determine a best-fit C/O ratio for WASP-12b of 0.59 ($\chi^{2}_{best}$/$N$=2.45, with a 68\% confidence interval of 0.54-0.95), suggesting that a high C/O ratio is not the explanation for WASP-12b's lack of atmospheric temperature inversion. If WASP-12b's C/O ratio really is super-solar and significantly different than its host star (0.48$\pm$0.08), this suggests that some other mechanism influenced the composition of the exoplanet during its formation/evolution. \"Oberg et al.\,(2011) note that the high C/O and substellar C/H reported by Madhusudhan et al.\,(2011) are only consistent with an atmosphere formed predominantly from gas accretion outside the water snowline. With our updated metallicity measurement, C/H in WASP-12 decreases to $\sim$ 5$\times 10^{-4}$, exactly in the middle of the C/H distribution spanned for the planet in Madhusudhan et al.\,(2011)'s best-fitting ($\chi^2<$7) models. Thus, by these models, the planet's C/H is just as likely to be substellar as super-stellar. More data, particularly around 3 $\mu$m (see Line et al.\,2013, Figure 1), can help further constrain WASP-12b's C/O ratio and enable a more meaningful comparison between planet and host star. We note that the very recent HST/WFC3 transit spectra of WASP-12b from 1.1-1.7 $\mu$m reported by Mandell et al.\,(2013) are fit equally well by oxygen- and carbon-rich models of Madhusudhan et al.\,(2012). \subsubsection{XO-1} The hot Jupiter XO-1b's (McCullough et al.\,2006) four \textit{Spitzer/IRAC} photometric secondary eclipse observations have been explained with a solar-composition, thermally-inverted model (Machalek et al.\,2008). However, it is also possible to fit the observations with a non-inverted (Tinetti et al.\,2010), potentially carbon-rich atmosphere model (Madhusudhan 2012), which may include disequilibrium chemistry like photochemistry and/or transport-induced quenching (Moses et al.\,2013). As the favored C/O$\geq$1 models are heavily dependent on the 5.8 $\mu$m photometric point, new observations are necessary to confirm the carbon-rich nature of XO-1b's atmosphere. Here we find in the XO-1 host star [Fe/H] = -0.11$\pm$0.06, with [C/H] = -0.19$\pm$0.04 and [O/H] = -0.09$\pm$0.05, resulting in C/O = 0.43$\pm$0.07. \subsubsection{TrES-2 and TrES-3} TrES-2b and TrES-3b were among the first transiting hot Jupiter exoplanets discovered (O'Donovan et al.\,2006; O'Donovan et al.\,2007). Both fall under the highly-irradiation ``pM'' class predicted to have temperature inversions in their upper atmospheres (Fortney et al.\,2008). Secondary eclipses of TrES-2b and TrES-3b were observed with the CFHT Wide-field Infrared Camera 2.15 $\mu$m filter, (Croll et al.\,2010a; 2010b), in \textit{Spitzer}/IRAC's four near-IR bands (O'Donovan et al.\,2010; Fressin et al.\,2010, respectively). Radiative transfer analyese of Line et al.\,(2014) use the CFHT and $\textit{Spitzer}$ data indicate a range of temperature-pressure profiles are too cool for TiO and VO to be in the gas phase, which suggests that these species do not cause thermal inversions. Line et al.\,(2014) find that the data provide minimal constraints on the abundances of H$_2$O, CO$_2$, CO, and CH$_4$, and thus TrES-2b's atmospheric C/O ratio. Their best fit ($\chi^{2}_{best}$/$N$=0.60) is 0.20, but their 68\% confidence interval spans 0.021-8.25. Interestingly, the C/O ratio of TrES-2 that we derive, 0.41$\pm$0.05, has the lowest error in our sample and also the second-lowest C/O ratio value in our sample. Hence, if TrES-2b accreted much of its gas from a reservoir similar in composition to its host star, and its atmosphere remained mostly isolated from its interior, it may also have a sub-solar atmospheric C/O ratio. For TrES-3b, radiative transfer analyses of the infrared photometry indicate that H$_2$O is well determined with an abundance near 10$^{-4}$, and CH$_4$ has an upper limit of $\sim$10$^{-6}$ (Line et al.\,2014). CO$_2$ shows a weak upper limit $\sim$10$^{-4}$, derived from the 2.1 $\mu$m CO$_2$ band wings within the $K$ band measurement, while CO is unconstrained due to the large uncertainty in the $Spitzer$ 4.5 $\mu$m data point, combined with the fact that no other molecular absorption features of CO are probed by the current data (Line et al.\,2014). Line et al.\,(2014) infer that C/O$>$1 in TrES-3b due to the relatively high-confidence limit on H$_2$O and the small upper limit on CH$_4$. However, the data provide no constraints on the CO abundance, which is expected to be the major carbon carrier in an atmosphere as hot as TrES-3b. Their best fit ($\chi^{2}_{best}$/$N$=0.067) C/O ratio for TrES-3b is 0.22, with a 68\% confidence interval of 0-0.97. The very recently published HST/WFC3 secondary eclipse observations of TrES-3b are poorly fit with a solar-composition model ($\chi^{2}/N=3.04$), whereas the WFC3 data plus the existing $\textit{Spitzer}$ photometry are more consistent ($\chi^{2}/N=0.75$) with an atmosphere model depleted in CO$_2$ and H$_2$O by a factor of 10 relative to a solar-composition model (Ranjan et al.\,2014). TrES-3's C/O ratio derived here, 0.29$\pm$0.09, also has a small error and is the lowest C/O ratio in the hot Jupiter host stars studied here. The large span in the planet's C/O ratio found by Line et al.\,(2014) is still too large to draw meaningful conclusions about the formation location and/or growth history of TrES-3b. However, if the degeneracy between CO and CO$_2$ absorption in the \textit{Spitzer} 4.5 $\mu$m data point is broken by, for instance, observations of the 2.6 $\mu$m or 15 $\mu$m CO$_2$ bending band or the 5 $\mu$m CO fundamental band by SOFIA/FLITECAM (McLean et al.\,2006) or SOFIA/FORCAST (Adams et al.\,2010), both the CO and CO$_2$ contributions could be better estimated and lead to a tighter C/O ratio constraint for TrES-3b. This system is intriguing and important for further investigation due to both our firmly sub-solar C/O ratio and the relatively metal-poor nature of the host star ([Fe/H]=-0.21$\pm$0.08, the lowest in our sample), which distinguishes TrES-3 from most other hot Jupiter hosts. \subsubsection{HD 149026} Observational constraints and extensive theoretical modeling indicate that the exoplanet HD 149026b has between 45-110 M$_{\oplus}$ of heavy elements in its core and surrounding envelope (Sato et al.\,2005; Fortney et al.\,2006; Ikoma et al.\,2006; Broeg \& Wuchterl 2007) , making the core of HD 149026b at least twice as massive as Saturn's, even though its radius is $\sim$0.86 R$_{\rm{Saturn}}$ (Triaud et al.\,2010) and its mass is $\sim$1.2 M$_{\rm{Saturn}}$ (Sato et al.\,2005). The massive core of HD 149026b challenges formation by traditional core accretion theory, and many modified formation scenarios have been suggested, including collision with an outer additional giant planet (Sato et al.\,2005; Ikoma et al.\,2006), accretion of planetesimals or smaller (super-Earth-sized) planets (Ikoma et al.\,2006; Broeg \& Wuchterl 2007; Anderson \& Adams 2012), or core accretion in a disk with $\times$2 the heavy element mass in the solar nebula (Dodson-Robinson \& Bodenheimer 2009). This latter explanation stems from the metal-rich nature of the star -- more massive/metal-rich disks form planets more readily (Ida \& Lin 2004ab) and metal-rich planets tend to be associated with metal-rich stars (Guillot et al.\,2006; Burrows et al.\,2007a; Miller \& Fortney 2011). Here we find [Fe/H]=0.26$\pm$0.09 for HD 149026, which is not as high as previous studies ([Fe/H]=0.36$\pm$0.05; Sato et al.\,2005), but still suggests that, overall, the initial metal abundance in the molecular cloud/disk was enhanced above solar. We measure [C/H]$=$0.26$\pm$0.08 and [O/H]$=$0.25$\pm$0.04, both enhanced above solar, resulting in a C/O ratio of 0.55$\pm$0.08, consistent with solar. Stevenson et al.\,(2012) find that the \textit{Spitzer} secondary eclipse observations of HD 149026b (at 3.6, 4.5, 5.8, 8.0, and 16 $\mu$m) can be fit using models with an atmosphere in chemical equilibrium and lacking a temperature inversion, with large amounts of CO and CO$_2$, and a metallicity $\times$30 solar (Fortney et al.\,2006). The retrieval results of Line et al.\,(2013) also indicate the atmosphere of HD 149026b has more CO and CO$_{2}$ than CH$_{4}$, which makes sense given the planet's high temperature ($\sim$1700 K) that favors formation of CO over CH$_{4}$ at solar abundances. There is also a peak in the Line et al. modeled composition probability distribution of the H$_2$O mixing ratio near $\sim$10$^{-5}$. Given this H$_2$O abundance, and the low abundance of CH$_{4}$, the C/O ratio of HD 149026b is likely $<$1, but is remains poorly constrained (0.55, with a $\chi^{2}_{best}$/$N$=0.23 fit and a 68\% confidence interval of 0.45-1.0; Line et al.\,2013). A better estimate of HD 149026b's C/O ratio as compared to the C/O ratio of its host star (0.55$\pm$0.08) may shed light on the planet's history and the origin of its massive core. This heavy-cored hot Jupiter system, with the host star carbon and oxygen abundances presented here, is a valuable test-bed for studying how massive planets form. \subsubsection{XO-2} The hot Jupiter XO-2b has a host star, XO-2N, with a binary companion, XO-2S, located $\sim$4600 AU away and not known to host a hot Jupiter-type planet (Burke et al.\,2007). The stars are of similar stellar type, meaning that the non-hosting companion can be used to check for effects of planet formation on the host star, e.g., stellar atmospheric pollution. XO-2b has been observed with with HST and \textit{Spitzer} (Machalek et al.\,2009; Crouzet et al.\,2012) as well as from the ground (Sing et al.\,2012 \& 2011; Griffith et al.\,2014). Griffith et al.\,(2014) find, with a comprehensive analysis of all existing data, that the water abundance that best matches most of the data is consistent with an atmosphere that has the same metallicity and C/O ratio as the host star in photochemical equilibrium. However there are outlying observations, so additional measurements and needed to understand the cause for the outliers and to investigate the carbon abundance in XO-2b. Teske et al.\,(2013b) derived the carbon and oxygen abundances of both binary components, and found [C/H]$=$ $+$0.26$\pm$0.11 in XO-2S versus $+$0.42$\pm$0.12 in XO-2N, and [O/H]$=$ $+$0.18$\pm$0.15 in XO-2S versus $+$0.34$\pm$0.16 in XO-2N. The stars are enhanced above solar in C and O, with XO-2N being slighly more carbon- and oxygen-rich. Their relative enhancements result in both having C/O=0.65$\pm$0.20. (Note that this value is slightly larger than that reported in Teske et al.\,2013b because the log$N$(O)$_{\odot}$ in this work is 8.66 versus 8.69 in Teske et al.\,2013b.) Both XO-2N and XO-2S fall exactly on our linear trends with [Fe/H] discussed in $\S$4.1 ([Fe/H]$_{XO-2N}$=0.39$\pm$0.14, [Fe/H]$_{XO-2N}$=0.28$\pm$0.14). The elevated-above-solar [C/H] and [O/H] values in the two stars are strong evidence that their parent molecular cloud was elevated in both carbon and oxygen. Given that their C/O ratios are identical, the key to understanding why XO-2N has a planet and XO-2S does not may lie in the exoplanet composition. \subsubsection{CoRoT-2} Of the planets around the host stars in our sample, CoRoT-2b is perhaps the most puzzling in terms of its atmospheric structure. Traditional solar composition, equilibrium chemistry models are unable to reproduce the unusual flux ratios from the three \textit{Spitzer} channel observations (it is missing 5.8 $\mu$m) of this very massive hot Jupiter (Alonso et al.\,2010; Gillon et al.\,2010; Deming et al.\,2011; Guillot \& Havel 2011). Despite its large mass, the planet has one of the greatest radius anomalies -- slower-contraction evolution models that explain the radius anomalies of other inflated planets cannot justify this case (Guillot \& Havel 2011). Furthermore, the host star is young (formed within 30-40 million years; Guillot \& Havel 2011), chromospherically active and the system has been suggested to be undergoing magnetic star-planet interactions due to the observed stellar spot oscillation period that is $\sim$10$\times$ the synodic period of the planet as seen by the rotating active longitudes (Lanza et al.\,2009). CoRoT-2b's emission data are difficult to interpret, largely because of the anomalously high 4.5/8.0 $\mu$m flux ratio. Excess CO mass loss has been suggested to enhance the 4.5 $\mu$m flux, as has some unknown absorber acting only below $\sim$ 5 $\mu$m (Deming et al.\,2011; Guillot \& Havel\,2011). Alternatively the low 8 $\mu$m flux may be caused by a high C/O ratio through absorption of CH$_4$, HCN, and C$_2$H$_2$ absorption (Madhusudhan\,2012). In addition, the lack of a 5.8 $\mu$m measurement leads to a poor constraint on the H$_2$O abundance, which strongly dictates the resulting C/O ratio. Wilkins et al.\,(2014) find that no single atmospheric model is able to reproduce of all the available CoRoT-2b data, including their new 1.1.7 $\mu$m HST/WFC3 spectra, the optical eclipse observed by \textit{CoRoT} (Alonson et al.\,2009; Snellen et al.\,2010) and the previously-modeled infrared photometry. More complex models with differing C/O ratios or varying opacity sources do not provide a fit more convincing that a one-component blackbody, which in itself still misses the ground-\textit{Spitzer} eclipse amplitudes by $\sim$1.8$\sigma$ (Wilkins et al.\,2014). Disequilibrium chemistry can significantly affect CoRoT-2b's atmospheric composition. For instance, for a high C/O ratio, H$_2$O is predicted to be enhanced above 10$^{-2}$ bar by $\sim$four orders of magnitude due to both transport-induced quenching and CO photochemistry in the upper atmosphere (Moses et al.\,2013). HCN and other C$_x$H$_x$ compounds may also result from the reaction of the leftover C with N or H$_2$ (Moses et al.\,2013). Disequilibrium chemistry models with 0.5$\times$solar metallicity, moderate mixing, and C/O=1.1 yield a significantly better match to the four CoRoT-2b infrared secondary eclipse observations, providing a $\chi^2/N$=1.3, versus the solar-composition models, which provide a $\chi^2/N$=7.2 (Moses et al.\,2013). The host star C/O ratio derived here, 0.47$\pm$0.09, is equal to C/O$_{\odot}$ within error, as is the overall metallicity of the host star, [Fe/H]=0.06$\pm$0.08. We note that the [C/H] value measured here for CoRoT-2 is based on only 2 carbon lines (5380\,{\AA} and 7113\,{\AA}), and the [O/H] value is based on only the O I triplet at $\sim$7775\,{\AA}; the other potential C and O lines were too weak to be reliably measured in our data. In addition, while the O I triplet NLTE corrections are nominally valid in this case because the T$_{\rm{eff}}$ of CoRoT-2 is $\geq$5400, implementing these corrections results in an [O/H] that is larger than the LTE case, the opposite direction of the corrections at near solar temperatures. However,[O/H]$_{\rm{LTE}}$=0.02, the same as [O/H]$_{\rm{NLTE}}$=0.06 within error (0.07 dex); if we adopt the [O/H]$_{\rm{LTE}}$ value, CoRoT-2's C/O ratio increases by only 0.05 dex, also within error. CoRoT-2b's C/O ratio, while still uncertain, could plausibly be $>$1. Several different scenarios could account for C/O$_{\rm{planet}}>$C/O$_{\rm{star}}$. CoRoT-2b could have accreted carbon-rich very hot gas from the inner disk regions ($\lesssim$0.1 AU). Alternatively, the planet could have accreted the majority of its gas from beyond the H$_2$O snow line (causing it to be oxygen-depleted), or accreted solid material depleted in oxygen (e.g., from a ``tar line'' inward of the snow line; Lodders 2004). Interestingly, to explain CoRoT-2b's inflated size and large mass, Guillot \& Havel\,(2010) propose that the CoRoT-2 system previously included multiple giant planets that collided within the last $\sim$20 million years to create the currently-observed CoRoT-2b. This scenario could result in a planet that differs significantly in composition from the original states of the impactors, potentially erasing the signatures of where/from what material in the disk the planet formed. CoRoT-2b is another candidate for which additional SOFIA observations at near- (2.6 $\mu$m, 6 $\mu$m) and mid-infrared ($>$20 $\mu$m) wavelengths can help better constrain the exoplanet C/O ratio and thus its formation history. \subsubsection{HD 189733} The hot Jupiter HD 189733b is one of the best-studied to date, with data spanning $\sim$0.3-24 $\mu$m (Barnes et al.\,2007; Grillmair et al.\,2007; 2008; Tinetti et al.\,2007; Knutson et al.\,2007; 2009; 2012; Redfield et al.\,2008; Charbonneau et al.\,2008; Beaulieu et al.\,2008; Pont et al.\,2008; D{\'e}sert et al.\,2009; Swain et al.\,2008 \& 2009a; Sing et al.\,2009 \& 2011; Agol et al.\,2010; Gibson et al.\,2012; Evans et al.\,2013; Birkby et al.\,2013). The first complete atmospheric study via statistical analysis with a systematic, wide parameter grid search (Madhusudhan \& Seager 2009) analyzed separately spectroscopic data from 5-14 $\mu$m (Grillmair et al.\,2008), photometric data at 3.6, 4.5, 5.8, 8, 16, and 24 $\mu$m (Charbonneau et al.\,2008), and spectrophotometric data from 1.65-2.4 $\mu$m (Swain et al.\,2009a). Madhusudhan \& Seager\,(2009) place constraints at the $\xi^2$=2 level (where $\xi^2$ is a proxy for the reduced $\chi^2$ using the \# of data points as $N$) on HD 189733b's atmospheric mixing ratios of H$_2$O, CH$_4$, and CO$_2$ using the spectrophotometric data, as it includes features of all of these molecules as well as CO. Their resulting C/O ratio range for HD 189733b is between 0.5 and 1. Subsequent analysis of all of the available infrared secondary eclipse measurements with the Bayesian optimal estimation retrieval scheme NEMESIS (Irwin et al.\,2008) placed constraints on molecular abundances ratios of HD 189733b's atmosphere, resulting in a best-estimate C/O ratio between 0.45 and 1 for $\xi^2<$0.5 and between 0.15 and 1 for $\xi^2<$2 (Lee et al.\,2012). However, these authors caution that the current secondary eclipse data are only able to constrain the thermal structure of HD 189733b at some pressure levels, and the mixing ratios of H$_2$O and CO$_2$ with large uncertainties ranging between 9-500$\times$10$^{-5}$ and 3-150$\times$10$^{-5}$ for $\xi^2<$0.5, respectively, due to the model degeneracies. The most significant degeneracy they find is between temperature and H$_2$O abundance at 300 mbar pressure. The H$_2$O abundance has the biggest influence on the overall shape of a hot Jupiter spectrum in thermochemical equilibrium. Moses et al.\,(2013) focuses on the H$_2$O mixing ratio constraint of $\sim$1$\times$10$^{-4}$ from Madhusudhan \& Seager\,(2009) in their exploration of disequilibrium chemistry using the combined data sets mentioned above. In both equilibrium and disequilibrum scenarios, for their nominal temperature profile and at solar metallicity, a very narrow range of C/O ratios around 0.88 provides the H$_2$O abundance constraint and a good fit to the observations. The recent retrieval analysis of Line et al.\,(2013), using the same wavelength coverage of data, also finds a best-fit C/O ratio of 0.85 ($\chi^{2}_{best}$/$N$=2.27 fit, with a 68\% confidence interval of 0.47-0.90). A carbon-enhanced atmosphere for HD 189733b is thus theoretically plausible and consistent with observations. Interestingly, we find the host star has C/O$=$0.90$\pm$0.15, matching well the best-fit C/O ratios derived for the planet's atmosphere. HD 189733 is the only star within this sample to have C/O$>$0.8; its C/O ratio spans 0.75-1.05 within 1$\sigma$ errors. Three additional stars in our sample have have C/O$>$0.8 within 1$\sigma$ errors. The derived T$_{\rm{eff}}$ of HD 189733 is $\leq$5400 K, therefore the triplet [O/H]$_{\rm{NLTE\,avg}}$ (0.125) is not included in the final average [O/H] reported here. Instead, the triplet [O/H]$_{\rm{LTE}}$ (0.01$\pm$0.14) and [O/H]$_{\rm{6300}}$ ($-$0.02$\pm$0.14) values are averaged. For stars as cool as HD 189733 there is evidence from studies of [O/H] in several open clusters that the canonical NLTE corrections are not appropriate -- [O/H]$_{\rm{LTE}}$ increases in lower-temperature stars in the same cluster, the opposite of what is predicted (Schuler et al.\,2006). If [O/H]$_{\rm{NLTE\,avg}}$ (0.125) is included in the average, the C/O ratio of HD 189733 is reduced to 0.82, and if [O/H]$_{\rm{NLTE\,avg}}$ replaces [O/H]$_{\rm{LTE}}$, the C/O ratio of HD 189733 is reduced to 0.79. The C/O ratio is reduced to 0.69 if [O/H]$_{\rm{NLTE\,avg}}$ is the sole oxygen abundance indicator. Alternatively, one may apply an empirical correction to [O/H]$_{\rm{LTE}}$ based on the temperature of HD 189733 and the observed cluster [O/H]$_{\rm{LTE}}$ anomaly (Schuler et al.\,2006), which amounts to $\sim$0.14 dex. This increases the resulting C/O ratio to 1.20. HD 189733's [C/H] is an outlier as compared to the rest of our sample, while its [O/H] is more consistent with the rest of the sample (Figure \ref{comparison_plot2}). Both measurements have some of the largest abundance errors of all the targets in our sample. Our reported [C/H]$=$0.22$\pm$0.13 for HD 189733 is in fact based on only one carbon line, 5380\,{\AA}, from the Keck/HIRES data, though we were able to measure two lines (5380 and 7111\,{\AA}) in the Subaru/HDS data, resulting in [C/H]$=$0.24$\pm$0.15. Including the 7113\,{\AA} C I line measurement from the Keck/HIRES data or the Subaru/HDS data increases the [C/H] from 0.22/0.24 to 0.34/0.33, and the reported C/O ratio to 1.16. Thus is appears that the C/O ratio of HD 189733 could be as low as$\sim$0.75, but is very likely $\gtrsim$0.80, as we report here. In order to match the desired H$_2$O mixing ratio, the C/O ratio of the exoplanet HD 189733b's atmosphere must shift to higher values when its metallicity is increased -- with an increase of 3$\times$solar in metallicity the C/O ratio reaches $\sim$0.96, compared to our derived C/O of 0.90$\pm$0.15. Alternatively, if the metallicity is sub-solar, the required C/O ratio decreases (Moses et al.\,2013). Unfortunately with present data the metallicity of HD 189733b's atmosphere is unknown. We find [Fe/H]=0.01$\pm$0.15 in the host star, providing at least a first-order constraint on the planetary atmospheric metallicity, but not a better constraint on its C/O ratio. However, we note that based on Moses et al.\,(2013)'s models, a change in the exoplanet's C/O ratio from 0.5 (solar) to 0.88 results in a change in the CH$_4$ abundance by $\sim$an order of magnitude, which should produce observable spectral signatures in the exoplanet's atmosphere. \subsubsection{Future of Direct Planet-Star Comparisons} The fact that HD 189733b is one of the most-studied hot Jupiters and yet still has a C/O ratio that can be anywhere from $\sim$0.5-1 indicates the difficulty and uncertainty in deriving exoplanet abundance ratios. Current limitations due larger to the paucity of data, which gives rise to degenerate solutions for transiting planet spectroscopy (e.g., Griffith 2013). However, as observational efforts continue to improve the quantity and quality of the measurements more precise C/O ratios will be possible. In addition, studies of transiting planets at high spectral resolution are becoming progressively refined (e.g., Snellen et al.\,2010; Birkby et al.\,2013; de Kok et al.\,2013; Brogi et al.\,2013) to the point that C/O ratio constraints are expected in the near future. Complementary sutides of younger, hotter planets are possible with spectroscopy of directly imaged planets. One such system, HR 8799, is particularly promising as it has four directly-imaged planets of similar luminosities, masses, and radii but different orbital distances and, surprisingly, maybe even different compositions (Barman et al.\,2011a; Currie et al.\,2011; Galicher et al.\, 2011; Marley et al.\,2012; Skemer et al.\,2012, 2013; Konopacky et al.\,2013). Recent directly-imaged, moderate-resolution ($R\sim$4000) spectra from $\sim$1.97-2.38 $\mu$m of HR 8799c show absorption of CO and H$_2$O but little to no CH$_4$ (mixing ratio $<$10$^{-5}$). $\chi^2$ minimization modeling of these data finds best-fit log$N$(C) and log$N$(O) values of 8.33 and 8.51, respectively, indicating HR 8799c is depleted in both C and O with respect to solar, and resulting in a C/O ratio of 0.65$^{+0.10}_{-0.05}$ (Konoapacky et al.\,2013). The host star is classified as both $\gamma$ Doradus and $\lambda$ Bootis, making stellar abundance analysis challenging, but one previous study of the star derives C/O=0.56$\pm$0.21 (Sadakane 2006). Determining the C/O ratios of the other planetary components in this system, and other multi-planet systems, may provide constraints on how the composition of the host star affects giant planet formation as a function of planet mass and orbital radius. | 14 | 3 | 1403.6891 |
1403 | 1403.6633_arXiv.txt | We show the effects of the perturbation caused by a passing by star on the Kuiper belt objects (KBOs) of our Solar System. The dynamics of the Kuiper belt (KB) is followed by direct $N$-body simulations. The sampling of the KB has been done with $N$ up to $131,062$, setting the KBOs on initially nearly circular orbits distributed in a ring of surface density $\Sigma \sim r^{-2}$. This modelization allowed us to investigate the secular evolution of the KB upon the encounter with the perturbing star. Actually, the encounter itself usually leads toward eccentricity and inclination distributions similar to observed ones, but tends also to excite the low-eccentricity population ($e\aplt 0.1$ around $a\sim 40$\,$\mathrm{AU}$ from the Sun), depleting this region of low eccentricities. The following long-term evolution shows a ``cooling" of the eccentricities repopulating the low-eccentricity area. In dependence on the assumed KBO mass spectrum and sampled number of bodies, this repopulation takes place in a time that goes from 0.5\,Myr to 100\,Myr. Due to the unavoidable limitation in the number of objects in our long-term simulations ($N \leq 16384$), we could not consider a detailed KBO mass spectrum, accounting for low mass objects, thus our present simulations are not reliable in constraining correlations among inclination distribution of the KBOs and other properties, such as their size distribution. However, our high precision long term simulations are a starting point for future larger studies on massively parallel computational platforms which will provide a deeper investigation of the secular evolution ($\sim 100\,$Myr) of the KB over its whole mass spectrum. | The Solar System is hedged by a ring composed of a huge number of small bodies: the Edgeworth-Kuiper belt \citep{Jewitt} (hereafter briefly called Kuiper belt, or KB). The Kuiper belt bears the signature it the early evolution of the Solar System, and a contains records of the end-state of the accretion processes occurred in that region. Therefore, the knowledge of the history of the Kuiper belt objects (KBOs) is relevant to be able to develop a full consensus of the formation of the Solar System. The majority of the KBOs are located between about 30\,$\mathrm{AU}$ and 90 $\mathrm{AU}$ from the Sun, but most are around the 2:3 resonance with Jupiter, at 39.5\,$\mathrm{AU}$ and at its 1:2 resonance, roughly around 48\,$\mathrm{AU}$. The total mass is estimated from $0.01$ to $0.1$ M$_\oplus$ \citep{Luu}. There are several, indirect, arguments suggesting that this is just a small fraction of its initial mass because most of it has been lost (see \cite{Kenyon}). The size distribution of the KBOs is, usually, assumed as a power law $dn/dR = A R^{-q}$, where $A$ and $q$ are constants. The $q$ exponent is estimated $\sim 4.0 \pm 0.5$ \citep{Ber,Fraser}. For a more detailed description of the Kuiper belt we refer, e.g., to \cite{Luu}. The KB has a bimodal inclination distribution resulting of two separate populations \citep{Brown}. The \emph{dynamically cold} population refers to objects moving on almost planar orbits with relatively low inclinations (up to about 10$^\circ$) respect to the ecliptic. On the other side, the dynamically hot population is characterized by highly inclined orbits (up to 40$^\circ$) with respect to the ecliptic. Note that these two populations are different from what we call, in this paper, the \emph{low-eccentricity} population, which are objects on nearly circular orbits (orbital eccentricities $< 0.1$) and the \emph{high-eccentricity} population (eccentricities $\gtrsim$ 0.1. In Fig. \ref{fig:oss} the eccentricities and inclinations are plotted as function of the semi-major axis for KBOs observed from the Minor Planet Center (MPC) \citep{mpc} which is the center of the Smithsonian Astrophysical Observatory (SAO) dedicated to tracking, monitoring, calculating and disseminating data from asteroids and comets. \begin{figure} \centering \subfigure{ \includegraphics[scale=0.3]{fig1a.pdf} } \subfigure{ \includegraphics[scale=0.3]{fig1b.pdf} } \caption{distribution of the eccentricities (top panel) and inclinations (bottom panel) as a function of semi-major axis of the observed KBOs; data are from the Minor Planet Center \citep{mpc}. The vertical lines highlight the main resonances.} \label{fig:oss} \end{figure} The KBOs have been sub-categorized in three groups: \begin{enumerate} \item classical KBOs ($42$ $< a < 49$ , $\langle e \rangle \simeq 0.09$, $\langle i \rangle \simeq 7^{\circ}$); \item scattered KBOs ($a > 30$ , $\langle e \rangle \simeq 0.49$, $\langle i \rangle \simeq 14^{\circ}$ ); \item main resonant KBOs : \begin{itemize} \item 4:3 resonance ($a \simeq36,4$, $\langle e \rangle \simeq 0.22$, $\langle i \rangle \simeq 8^{\circ}$); \item 3:2 resonance, Plutino's ($a \simeq39.4$, $\langle e \rangle \simeq 0.36$,$\langle i \rangle \simeq 13^{\circ}$); \item 2:1 resonance ($a \simeq47.8$, $\langle e \rangle \simeq 0.14$, $\langle i \rangle \simeq 10^{\circ}$)., \end{itemize} \end{enumerate} where semimajor axes, $a$, are in $\mathrm{AU}$. These sub-populations have been explain through a phase of planet migration and a phase of clearing of the environment during the evolution of the early Solar System \cite{Mal1, Mal2}. In the latter phase the resonance population was formed by sweeping resonance capture in which the Jovian planets withstand considerable orbital migration as a result of encounters with residual planetesimals. While Neptune moved outwards, a small body like Pluto in an initially circular orbit could have been captured into the 3:2 resonance. The high orbital eccentricity would subsequently be induced by repeated orbital crossings with Neptune. Many others studies have attempted to better understand the properties of the KBOs. \cite{Gomes} investigated how the outward migration of Neptune, as proposed by \cite{Mal1, Mal2}, could have scattered objects from $25$ $\mathrm{AU}$ onto high-$i$ orbits leading to the current classical Kuiper belt region. He concluded that the high-$i$ population was formed closer to the Sun and brought into the classical Kuiper belt during planetary migration, whereas the cold population represents a primordial, relatively undisturbed population. This also led to the speculation that other mechanisms, such as planetary migration, have been the cause of the correlation between inclinations and colors in the classical Kuiper belt rather than environmental effects like the collisions among the KBOs (see \cite{Dore}). Detailed discussions about the correlation of the inclination with the color, size and binary of the KBOs are given by \cite{Levi1, Bru, Noll,Volk}. More recently a model, called the Nice model, has been proposed \citep{Levi}, which argues that the giant planets migrated from an initial compact configuration into their present orbits, long after the dissipation of the initial protoplanetary gas disk. The Nice model seems to provide a acceptable explanation for the formation of the classical and scattered populations, and for the correlation between inclinations and colors (for more detail see \cite{Levi}). The Nice model, however, predicts a higher eccentricities in classical KBO orbits than is observed. An interaction between a passing field star and the the Solar System could also be responsible for some of the orbital families observed in the KBO, which is the main topic of this paper. | We investigated the effect of an encounter between a passing star on the morphology of the Kuiper belt, and its subsequent long-term evolution. Using the current morphology of the KB we constrained the parameter of the incoming star. The orbit of the encountering star, the planets and those of the KBOs were integrated directly, as was the subsequent evolution of the internal dynamics of the KB and planets. The initial conditions for the Solar System ware taken from \cite{Ito}, and the Kuiper belt objects were distributed in a flat disk between 42 and 90 $\mathrm{AU}$ in the plane of the Ecliptic, and with a power law density distribution with exponent $-2$. The total mass of Kuiper belt ranged between 1 and 30 M$_\oplus$, and the mass of the incoming star was chosen to be 0.5\,\MSun\, 1.0 and 2.0\,\MSun. We compared the morphology of the KB directly after the encounter with the passing star, and after a secular evolution of up to 0.9\,Myr. The best results, directly after the encounter, are obtained when the incoming star approached the ecliptic plane with an impact parameter of 170-220 $\mathrm{AU}$ and an inclination above the Ecliptic of $60^\circ$ to $120^\circ$. The lower (best) values of both $b$ and $\theta$ are for the 0.5\,\MSun, encountering star whereas the upper values correspond to the 2.0\,\MSun\, intruder. We summarize these results in Tab. \ref{tab:conc}. In Fig. \ref{fig:Mstar_vs_d} we present the impact parameter and the angle $\theta$ of the incoming star as a function of its mass. A correlation between these parameters is evident and shows a degeneration in the parameters space. In fact, using different parameters is possible to reproduce an encounter with the same strength and find similar proprieties in the final KBOs distributions. During this encounter about 13\% of the Kuiper belt is lost from the Solar system. Actually, results do not show a depletion of the original flat distribution up to $\sim 99$ $\%$ as suggested by the observed total mass of the KB, evaluated in the range $0.3-0.1$ M$_\oplus$, and the mass estimation from Solar system formation model, $30-10$ M$_\oplus$, \citep{Luu} and match the eccentricity and inclination distributions with the observation at the same time. On the other hand, a better coverage of the initial conditions of the incoming star can very likely enhance the possibility of finding an {\it intermediate} case, where a strong depletion can be compatible with the observed distributions in eccentricities and inclinations. The morphology of the high-eccentricity and scattered population of the KB are well represented directly after the encounter. The low-eccentricity population, around $\sim 40$ $\mathrm{AU}$ and with eccentricities $\aplt 0.1$ is almost completely absent directly after the encounter. This mismatch in the morphology can be resolved by taking the secular evolution of the Kuiper belt into account. The low-eccentricity population is reinstated within a million years. Our models did not show any particular correlation between the inclinations distributions and the mass of the KBOs. However, due to the limited number of objects in our simulations, we could run only almost single-mass particle simulations. For example in our best model the gap in mass between the two population is only a factor 5 and the ratio between the radius is 1.7. Due to this limitation it is not possible to constrain any significant correlations among inclination and other properties, such as the size distribution of the KBOs and the number of KBO binaries. In conclusion, the sampling limiting our model and the relatively short time-scale of our simulations cannot give reliable results on that (actually, our finest simulation involved 16384 KBOs and was carried up to 0.9 Myear). We expect that a more sophisticated investigation of the long-term ($\sim 100\,$Myr) evolution of the KB, with a proper population over the whole KBO mass spectrum will show the ``relaxation" of the eccentricities to low values as it happens in the case of mass monodisperse-particle simulations. Such detailed studies would thus provide important information about the final distribution of the KBOs, which will allow a complete comparison with observable such as the size-inclination relation, but unfortunately it is hard to achieve at the moment without the access to a very large GPUs cluster. While the secular evolution repopulated the low-eccentricity population, it triggered the further KB causing the depletion of the resonance population, which was initiated by the passing star. This loss of the resonant population can be due to the insufficient sampling of the KB in our simulations. Alternatively, the early migration of the planets is driving the repopulation of the resonant families \citep{Mal2, Ida1}. Such planetary reordering would be a natural consequence of the Nice model \citep{Levi}. Moreover, the resonances, that we have suddenly after the passage of the fly-by star, do not show an eccentricity and inclination distributions compatible with the observations. \begin{table} \centering \begin{tabular}{|c|c|c|c|} \hline M & 0.5 & 1.0 & 2.0 \tabularnewline \hline $b$ & $170$ & $200$ & $220$ \tabularnewline $v_{\infty}$& $3$ & $3$ & $3$ \tabularnewline $\theta$ & $60$ & $90$ & $120$ \tabularnewline \hline \hline \end{tabular} \caption{The optimal encounter parameters ($b$ and $\theta$) obtained for a star with mass $M$ (in solar masses) approaching with velocity at infinity of 3 km/s the Solar System (model B). The units of $b$ and $\theta$ are those adopted in this paper.} \label{tab:conc} \end{table} \begin{figure} \centering \includegraphics[width=0.5\textwidth]{fig14.pdf} \includegraphics[width=0.5\textwidth]{fig15.pdf} \caption{The impact parameter $b$ (top) and inclination $\theta$ (bottom) as a function of the mass of the incoming star for our most favorite model (see Tab. \ref{tab:conc}). The dashed line gives a fit to the tree points to indicate the trend, which follows $b = 170 + 45(M-0.5)$ for the impact parameter and $\theta = 65 + 40(M_\star-0.5)$ for the inclination.} \label{fig:Mstar_vs_d} \end{figure} | 14 | 3 | 1403.6633 |
1403 | 1403.4308_arXiv.txt | We present deep VERITAS observations of the blazar PKS 1424+240, along with contemporaneous \textit{Fermi} Large Area Telescope, \textit{Swift} X-ray Telescope and \textit{Swift} UV Optical Telescope data between 2009 February 19 and 2013 June 8. This blazar resides at a redshift of $z\ge0.6035$, displaying a significantly attenuated gamma-ray flux above 100~GeV due to photon absorption via pair-production with the extragalactic background light. We present more than 100 hours of VERITAS observations from three years, a multiwavelength light curve and the contemporaneous spectral energy distributions. The source shows a higher flux of (2.1$\pm0.3$)$\times10^{-7}$ ph m$^{-2}$s$^{-1}$ above 120 GeV in 2009 and 2011 as compared to the flux measured in 2013, corresponding to (1.02$\pm0.08$)$\times10^{-7}$ ph m$^{-2}$s$^{-1}$ above 120 GeV. The measured differential very high energy (VHE; $E\ge100$ GeV) spectral indices are $\Gamma=$3.8$\pm$0.3, 4.3$\pm$0.6 and 4.5$\pm$0.2 in 2009, 2011 and 2013, respectively. No significant spectral change across the observation epochs is detected. We find no evidence for variability at gamma-ray opacities of greater than $\tau=2$, where it is postulated that any variability would be small and occur on longer than year timescales if hadronic cosmic-ray interactions with extragalactic photon fields provide a secondary VHE photon flux. The data cannot rule out such variability due to low statistics. | PKS\,1424+240 (VER\,J1427+237) is a distant very high energy (VHE; $E\ge100$ GeV) blazar at $z\ge0.6035$ \citep{furniss1424}. At this \textit{minimum} distance, the intrinsic VHE emission is expected to be significantly absorbed by the extragalactic background light (EBL) via pair-production, $\gamma + \gamma \rightarrow e^{+} + e^{-}$ \citep{nikishov}. The absorption of VHE gamma rays by the EBL can be estimated using the model-dependent gamma-ray opacity, $\tau(E,z)$. The source flux, $F_{\rm int}$, can be estimated from the observed flux, $F_{\rm obs}$, using the relation $F_{\rm int} = F_{\rm obs}\times e^{\tau(E,z)}$. The EBL cannot be directly measured due to foreground sources. The modification of distant VHE blazar spectra has been used to estimate the spectral properties of the EBL \citep{aharonian2006,albert2008}, providing photon density upper limits consistent with the observational lower limits set by galaxy counts \citep{werner}. Recent work has indicated that the EBL density is closer to the lower limits than the upper limits \citep{HESSEBL,hornsmeyer,fermiEBL}. The distance to PKS\,1424+240 makes the source ideal for studying extragalactic VHE photon propagation. % The high-energy spectral energy distribution (SED) measured in initial observations by VERITAS and the \textit{Fermi} Large Area Telescope (LAT) \citep{acciari1424} is investigated in \cite{furniss1424}, showing an absorption-corrected spectrum suggestive of VHE spectral hardening, though not beyond the conservative $\Gamma=1.5$ spectral limitation (where $dN/dE \propto E^{-\Gamma}$) described in, e.g., \cite{aharonian2006}. In an effort to understand the gamma-ray emission from PKS\,1424+240, we analyze deeper observations by VERITAS and LAT, including more than four times the exposure in \cite{acciari1424} and \cite{furniss1424}. In order to minimize hardening introduced from EBL absorption corrections, we explore the gamma-ray observations using the low-density ``fixed" model from \cite{gilmore2012}. This model, also providing compatible fits to LAT data in \cite{fermiEBL}, is comparable with that of \cite{franceschini} used in \cite{HESSEBL}, and provides similar absorption-corrections as compared to other EBL models, e.g. \cite{kneiske,dominguez,finke}. Luminosities calculated in this work use a H$_0$ = 100 $h$ km s$^{-1}$Mpc$^{-1}$ where $h $= 0.7. | The blazar PKS\,1424+240 resides at $z\ge0.6035$, with a VHE flux that is significantly attenuated by the EBL. Discovery observations of this source by VERITAS have shown a marginal indication of spectral hardening at the highest energies, after correction by the EBL \citep{furniss1424}. While a similar effect is seen in the deep observations obtained in 2013, the significance of the effect remains marginal because of the lower overall flux level during this epoch. In both epochs the data are consistent with a simple power law, even after correction for absorption by the EBL. If the indication of spectral hardening could be confirmed, one possible explanation would be an over-estimation of the EBL density, although the results shown use one of the lowest density EBL models currently available, which approaches the galaxy count lower limits of the EBL density at $z\sim0$. % The possible spectral hardening is of great interest, because if it is not from over-estimation of the EBL, there are a number of physical mechanisms which can produce hardening with increasing energy. Second-order synchrotron self Compton emission, pair-cascades initiated by pion decay in hadronic emission scenarios \citep{boettcher} or internal photon-photon absorption \citep{aharonian2008} can produce hard components at high energy. There are also scenarios that describe spectral hardening as arising from hadronic cosmic-ray line-of-sight interactions with the cosmic microwave background and EBL. These processes can produce secondary gamma rays close to the observer, hardening the observed VHE spectrum \citep{essey2010a,essey2010b,essey2011,essey2012,murase2012,razzaque2012,prosekin2012,aharonian2013,zheng2013,kalashev2013,inoue2013}. This component is expected to become dominant at high energies where the EBL opacity is greater than $\sim$2 and is not expected to vary on timescales shorter than about a year. The VERITAS observations above 310 GeV (where $\tau=2$ according to \citealt{gilmore2012} and \citealt{kneiske}) do not show significant variability between 2009 and 2013, nor can they strongly exclude it. More exotic theories, involving Lorentz invariance violation \citep{LIV} or axion-like particles (ALPs), might also produce spectral hardening at high energies, e.g. \cite{sanchezconde}. The blazar can be categorized as an high-synchrotron-peaked (HSP) BL Lac, with a synchrotron peak above 10$^{15}$ Hz \citep{abdoSED} and an isotropic luminosity above 400 GeV of 1.03$\times10^{44}$ erg s$^{-1}$. At $z\ge0.6035$, it is apparent that PKS\,1424+240 represents a powerful tool for studying intrinsic emission mechanism(s) within blazar jets, extragalactic cosmic-ray propagation and the propagation of VHE photons across extragalactic space. Future studies will benefit from additional VHE observations as well as from any additional information that will be obtained about the redshift, e.g. from HST/STIS UV observations. | 14 | 3 | 1403.4308 |
1403 | 1403.4764_arXiv.txt | We detect the second known $\lambda$ Bootis star (HD 54272) which exhibits $\gamma$ Doradus type pulsations. The star was formerly misidentified as a RR Lyrae variable. The $\lambda$ Bootis stars are a small group (only 2\%) of late B to early F-type, Population I stars which show moderate to extreme (up to a factor 100) surface underabundances of most Fe-peak elements and solar abundances of lighter elements (C, N, O, and S). The photometric data from the Wide Angle Search for Planets (WASP) and All Sky Automated Survey (ASAS) projects were analysed. They have an overlapping time base of 1566~d and 2545~d, respectively. Six statistically significant peaks were identified ($f_{1}=1.410116$\,d$^{-1}$, $f_{2}=1.283986$\,d$^{-1}$, $f_{3}=1.293210$\,d$^{-1}$, $f_{4}=1.536662$\,d$^{-1}$, $f_{5}=1.15722$\,d$^{-1}$ and $f_{6}=0.22657$\,d$^{-1}$). The spacing between $f_{1}$ and $f_{2}$, $f_{1}$ and $f_{4}$, $f_{5}$ and $f_{2}$ is almost identical. Since the daily aliasing is very strong, the interpretation of frequency spectra is somewhat ambiguous. From spectroscopic data, we deduce a high rotational velocity (250$\pm$25\,km\,s$^{-1}$) and a metal deficiency of about $-$0.8 to $-$1.1\,dex compared to the Sun. A comparison with the similar star, HR 8799, results in analogous pulsational characteristics but widely different astrophysical parameters. Since both are $\lambda$ Bootis type stars, the main mechanism of this phenomenon, selective accretion, may severely influence $\gamma$ Doradus type pulsations. | The group of classical $\lambda$ Bootis stars comprises late B to early F-type, Population I stars, with moderate to extreme (up to a factor of 100) surface underabundances of most Fe-peak elements and solar abundances of lighter elements (C, N, O, and S). They are rare, with a maximum of about 2\% of all objects in the relevant spectral domain, between the zero- and terminal-age main-sequence (ZAMS and TAMS), are found to be such objects (\citealt{Pau02b}). \citet{Mic86} suggested that the peculiar chemical abundances on the stellar surfaces are due to selective accretion of circumstellar (CS) material. Due to gravitational settling and radiative acceleration, it is then mixed in the shallow convection zone of the star. This explains why the anomalous abundance pattern is similar to that found in the gas phase of the interstellar medium (ISM) in which refractory elements like iron and silicon have condensed into dust grains. Later on, \citet{Kam02} and \citet{Mar09} developed a model which describes the interaction of the star with its local ISM and/or CS environment. As a result, different levels of underabundance are produced by different amounts of accreted material relative to the photospheric mass. The small fraction of this star group on the main-sequence (MS) is explained by the low probability of a star-cloud interaction and by the effects of meridional circulation, which dissolves any accretion pattern a few million years after the accretion has stopped. The hot end of this model is due to significant stellar winds for stars with $T_\mathrm{eff}$\,$>$\,12\,000\,K whereas the cool end, at about 6500\,K, is defined by convection which prevents the accreted material manifesting at the stellar surface. Strong support for the selective accretion scenario has been given by \citet{Fol12} who found that half of their sample of Herbig Ae/Be stars exhibit the characteristic $\lambda$ Bootis type abundance pattern. We know that the density of CS material around Herbig Ae/Be stars is very high, perfectly suited as the source for accretion. Almost all $\lambda$ Bootis stars are located within the classical $\delta$ Scuti/$\gamma$ Doradus instability strip. \citet{Pau02a} presented a detailed analysis of their pulsational behaviour. They concluded that at least 70\% of the group members inside the classical instability strip pulsate, and they do so with high overtone p-modes (Q\,$<$\,0.020\,d). In this paper, we present the newly detected $\gamma$ Doradus pulsational characteristics of the classical $\lambda$ Bootis star, HD 54272. Interesting enough, it was misidentified as a RR Lyrae type star by \citet{szczygiel2007}. We performed a detailed time series analysis of the Wide Angle Search for Planets (WASP) and All Sky Automated Survey (ASAS) photometric data. Four statistically significant frequencies and their combinations were detected. Our findings are discussed in comparison with the first hybrid $\lambda$ Bootis/$\gamma$ Doradus star, HR 8799. The latter is a very interesting young object hosting at least four planets and a massive dusty debris disk \citep{Esp13}. HD 54272, on the other hand, is a fast rotating star almost at the TAMS with no detected IR excess. From spectroscopy, we deduce a metallicity of $-$0.8 to $-$1.1\,dex compared to the Sun. This fact together with the high projected rotational velocities, make this object an interesting and important test case for models dealing not only with $\gamma$ Doradus pulsation, but also with selective accretion in the presence of meridional circulation. \begin{table} \centering \begin{minipage}{55mm} \caption{The astrophysical parameters of HD 54272 \citep[][this work]{Pau02b} and HR 8799 \citep{Gra99,Bai12}. The errors in the final digits of the corresponding quantity are given in parenthesis.} \label{stars_param} \centering \begin{tabular}{lcc} \hline\hline \multicolumn{1}{l}{\multirow{2}{*}{Quantity}} & HD 54272 & HR 8799 \\ & \multicolumn{2}{c}{values} \\ \hline \Teff [K] & 7010(217) & 7430(75) \\ $[$M/H$]$ & $-$0.8 to $-$1.1 & $-$0.47(10) \\ \logg & 3.83(10) & 4.35(5) \\ $M_{\mathrm{V}}$ [mag] & 2.33(30) & 2.98(8) \\ $M$ [M$_{\sun}$] & 1.69(19) & 1.47(30) \\ $R$ [R$_{\sun}$] & 2.2(3) & 1.34(5) \\ $L$ [L$_{\sun}$] & 1.01(12) & 0.69(5) \\ \logt & 9.12 & $<$\,8.00 \\ $v\,\sin\,i$ [km\,s$^{-1}$] & 250(25) & 38(2) \\ \hline \\ \end{tabular} \end{minipage} \end{table} \begin{figure} \begin{center} \includegraphics[width=85mm,clip]{figure1.eps} \caption{The classification resolution spectrum of HD 54272 together with the two standard stars HD 23194 (A5 V) and HD 113139 (F2 V). The latter were taken from \citet{Gra03} and have the same resolution as our spectrum.} \label{spec_hd54272} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=85mm,clip]{figure2.eps} \caption{The observed and synthetic (\Teff\,=\,7000\,K and \logg\,=\,3.8) spectra for HD 54272.} \label{synthetic} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=85mm,clip]{figure3.eps} \caption{The observed and two synthetic spectra showing the effect of metallicity (upper panel). The generated `rotating' model (see text) is not very different to the non-rotating one (lower panel).} \label{synthetic_two} \end{center} \end{figure} | We analysed the time series of the WASP and ASAS projects in order to shed more light on the pulsational behaviour of HD 54272. This object was classified as $\lambda$ Bootis star and RR Lyrae variable. A combination which, a-priori, exclude each other. Our detailed analysis established this object as a $\gamma$ Doradus pulsator with six detected frequencies, which could be possibly biased by one day aliases. The amplitudes of detected peaks are between 5 and 21\,mmag. The three frequencies with the highest amplitudes were detected in both data sets, independently. This example is also important for other $\gamma$ Doradus stars which might be misidentified as RR Lyrae variables. From spectroscopic data, we deduce a classification of `kA5hF2mA5V $\lambda$ Boo' for HD 54272. A comparison with synthetic spectra yields a high rotational velocity (250$\pm$25\,km\,s$^{-1}$) and a metal deficiency of $[$M/H$]$ of about $-$0.8 to $-$1.1\,dex compared to the Sun. It is located very close to the TAMS. These facts make HD 54272 an important test case for several different models including pulsation, rotation, diffusion, and selective accretion. Our results are compared with those from the first detected star which shows $\lambda$ Bootis and $\gamma$ Doradus characteristics: HR 8799. Although the pulsational behaviours are very similar, the evolutionary statuses are quite different. HR 8799 is a very young object hosting at least four planets and a massive dusty debris disk. If we accept that selective accretion plays a key role for the $\lambda$ Bootis phenomenon, it seems to have a significant effect on $\gamma$ Doradus type pulsation. This could be concluded from the fact that these two stars have widely different evolutionary statuses and astrophysical parameters, but a strikingly similar pulsational characteristics. The detection of further members of the $\lambda$ Bootis group showing $\gamma$ Doradus type pulsation would put further constraints on models explaining and describing these phenomena. | 14 | 3 | 1403.4764 |
1403 | 1403.3691_arXiv.txt | {High-density environments are crucial places for studying the link between hierarchical structure formation and stellar mass growth in galaxies. In this work, we characterise a massive proto-cluster at $z=2.895$ that we found in the COSMOS field using the spectroscopic sample of the VIMOS Ultra-Deep Survey (VUDS). This is one of the rare structures at $z\sim 3$ not identified around an active galactic nucleus (AGN) or a radio galaxy, thus it represents an ideal laboratory for investigating the formation of galaxies in dense environments. The structure comprises 12 galaxies with secure spectroscopic redshift in an area of $\sim7'\times8'$, in a total $z$ range of $\Delta z = 0.016$. The measured galaxy number overdensity is $\delta_g =12 \pm 2$. This overdensity has a total mass of $M\sim8.1\times10^{14}M_{\odot}$ in a volume of $13\times15\times17$ Mpc$^3$. Simulations indicate that such an overdensity at $z\sim2.9$ is a proto-cluster, which will collapse in a cluster of total mass $M_{z=0}\sim2.5 \times 10^{15} M_{\odot}$ at $z=0$, i.e. a massive cluster in the local Universe. We analysed the properties of the galaxies within the overdensity, and we compared them with a control sample at the same redshift but outside the overdensity. We could not find any statistically significant difference between the properties (stellar mass, star formation rate, specific star formation rate, NUV-r and r-K colours) of the galaxies inside and outside the overdensity, but this result might be due to the lack of statistics or possibly to the specific galaxy population sampled by VUDS, which could be less affected by environment than the other populations not probed by the survey. The stacked spectrum of galaxies in the background of the overdensity shows a significant absorption feature at the wavelength of Ly$\alpha$ redshifted at $z=2.895$ ($\lambda=4736$\AA), with a rest frame equivalent width ($EW$) of $4\pm 1.4$\AA. Stacking only background galaxies without intervening sources at $z\sim2.9$ along their line of sight, we find that this absorption feature has a rest frame EW of $10.8\pm3.7$\AA, with a detection S/N of $\sim4$. We verify that this measurement is not likely to be due to noise fluctuations. These EW values imply a high column density (N(HI)$\sim3-20\times10^{19}$cm$^{-2}$), consistent with a scenario where such absorption is due to intervening cold streams of gas that are falling into the halo potential wells of the proto-cluster galaxies. Nevertheless, we cannot rule out the hypothesis that this absorption line is related to the diffuse gas within the overdensity.} | The detection and study of (proto) clusters at high redshift is important input for cosmological models, and these high-density environments are crucial places for studying the link between hierarchical structure formation and stellar mass growth in galaxies at early times. The earlier the epoch when an overdensity is detected, the more powerful the constraints on models of galaxy formation and evolution because of the shorter cosmic time over which physical processes have been able to work. Specifically, in high-redshift ($z\gtrsim1.5-2$) (proto) clusters the study of how environment affects the formation and evolution of galaxies is particularly effective, because galaxies had their peak of star formation at such redshifts \citep{madau1996,cucciati12sfrd}. However, the sample of high-redshift ($z\gtrsim 1.5$) structures detected so far is still limited, and it is very heterogeneous, spanning from relaxed to unrelaxed systems. Many detection techniques have been used that are based on different (and sometimes apparently contradicting) assumptions. For instance, the evolution of galaxies in clusters appears to be accelerated relative to low-density regions (e.g. \citealp{steidel05}). As a result, while the average SFR of a galaxy decreases with increasing local galaxy density in the low-redshift Universe, this trend should reverse at earlier times, with the SFR increasing with increasing galaxy density \citep{cucciati2006,elbaz2007}. Indeed, some (proto-) clusters have been identified at high redshift as overdensities of star-forming galaxies \citep{capak11}, such as Ly$\alpha$ emitters \citep{steidel00,Ouchi03,Ouchi05,lemaux09} and H$\alpha$ emitters \citep{hatch11b}. At the same time, in some high-$z$ overdensities an excess of massive red galaxies has also been observed (e.g. \citealp{kodama07,spitler12}), and other dense structures have been identified via a red-sequence method (e.g. \citealp{andreon09}) or via an excess of IR luminous galaxies (e.g. \citealp{gobat11,stanford12}) or LBGs (e.g. \citealp{toshikawa12}). High-$z$ overdensities have also been identified by using other observational signatures, for instance with the Sunyaev-Zeldovich effect \citep{sunyaev_zeldovich1972,sunyaev_zeldovich1980} as in \cite{foley11_SZ}, or searching for diffuse X-ray emission (e.g. \citealp{fassbender11}), or looking for photometric redshift overdensities in deep multi-band surveys \citep{castellano07,salimbeni09,scoville13,chiang14_cosmos}. Moreover, assuming a synchronised growth of galaxies with that of their super-massive black holes, high-redshift proto-structures have been searched for around AGNs (e.g. \citealp{miley04}) and radio galaxies (e.g. \citealp{pentericci00,matsuda09,galametz12}), even if an excess of galaxies around high-$z$ QSOs has not always been found (see e.g. \citealp{decarli12}). However, this approach could introduce unknown selection effects, for example those due to the influence of powerful radio galaxies on the surrounding environment. The study of proto-structures selected only on the basis of the redshift distribution of its members is more likely to offer an unbiased view of high-density environments at high redshift and allow a comparison with the habitat of radio galaxies and quasars. Nevertheless, it is necessary to obtain spectroscopic redshifts of member galaxies, which is a costly observational task at high redshifts. Spectroscopic surveys conducted with visible wavelength spectrographs will observe the UV rest frame light of galaxies at redshifts $z>2$, and therefore be mostly sensitive to star-forming galaxies. Although the sample of (proto) clusters at $z>1.5$ is increasing in number, it is a heterogeneous data set. This inhomogeneity prevents using it to assess the abundance of clusters at such redshifts, which could be used to constrain cosmological parameters (e.g. \citealp{borgani1999,ettori09}). \cite{chiang13_sim} have recently made an attempt to find a common parameter to group and analyse the known overdensities at high $z$. They used simulations to study the probability of given overdensities at $z=2-5$ to collapse in bound clusters at $z=0$, and, in case of collapse, the mass at $z=0$ ($M_{z=0}$) of such clusters. They also give prescriptions for computing $M_{z=0}$ using the overdensity of the proto-cluster within a given volume. Following their own prescriptions, in a second work \citep{chiang14_cosmos} they perform a homogeneous search for overdensities using the photometric redshifts in the COSMOS field. We come back to their analysis in the following sections. The discovery and study of an overdensity at high $z$ also naturally addresses how a dense environment affects galaxy formation and evolution. Galaxies can build their stellar masses via abrupt processes like mergers, which in some cases produce an increase in mass up to a factor of two or so, or via more continuous processes based on in-situ star formation. At the same time, other physical processes are likely at work to quench star formation (such as AGN and SNe feedback), and some of these processes are particularly effective in high-density environments, where the gas reservoirs in galaxies can be stripped during interactions with the intra-cluster medium (ICM). The relative role of all these processes as a function of cosmic time is still a matter of debate. In recent years, many observational studies have focused on analysing the merger rate. If at $z<1$ the evolution of merger rate is quite well constrained for both major and minor mergers (i.e. with a luminosity/mass ratio greater or less than $\sim1/4$, see e.g. \citealp{deravel09} and \citealp{lopez_sanjuan11_vvds}), at $z>1$ observational results still show a large scatter (see e.g. \citealp{lopez_sanjuan13_massiv} and \citealp{tasca14_merger} for the most recent studies). On the side of stellar mass growth via smooth star formation, some theoretical models support a scenario where massive ($M_{baryon} \sim 10^{11} M_{\odot}$ ) galaxies at $z \sim 2 - 3$ are efficiently fed by narrow, cold (e.g. $T\sim 10^4$ K), intense, partly clumpy, gaseous streams that penetrate the shock-heated halo gas into the inner galaxy with rates of the order of 100 $M_{\odot} {\rm yr}^{-1}$ . These streams can grow a dense, unstable, turbulent disc with a bulge and trigger rapid star formation (e.g. \citealp{keres05,dekel09}). Observational evidence of gas accretion is still limited \citep{giavalisco11,bouche13}, and further studies are needed to support this scenario. Simulations \citep{kimm11} show that the covering fraction of dense cold gas is larger in more massive haloes, suggesting that the best environment for testing the cold flow accretion scenario are high-redshift over-dense regions. In this paper, we present the discovery of an overdensity at $z\sim2.9$ in the COSMOS field, detected in the deep spectroscopic survey VUDS (VIMOS Ultra-Deep Survey). In Sect.~\ref{data} we describe our data. In Sect.~\ref{overdensity} we describe the overdensity and compute the total mass that it comprises, and also its possible evolution to $z=0$. Section \ref{background} shows the search for diffuse cold gas in the overdensity, as inferred by absorption lines in the spectra of background galaxies. In Sect.~\ref{galaxies} we analyse the properties of the galaxies in the overdensity and contrast them to a sample of galaxies outside the structure at a similar redshift. Finally, in Sect.~\ref{discussion} we discuss our results and summarise them in Sect.\ref{summary}. We adopt a flat $\Lambda$CDM cosmology with $\Omega_m=0.27$, $\Omega_{\Lambda}=0.73$, and $H_0=70$ km~s$^{-1}$~Mpc$^{-1}$. Magnitudes are expressed in the AB system. | \label{summary} In this work, we have characterised a massive proto-cluster at $z=2.895$ that we found in the COSMOS field using the spectroscopic sample of the VUDS survey. Our results can be summarised as follows: \begin{itemize} \item[-] The overdensity comprises 12 galaxies with secure spectroscopic redshift in an area of $\sim7'\times8'$, in a total redshift range of $\Delta z = 0.016$. The measured galaxy overdensity is $\delta_g =12 \pm 2$. According to simulations \citep{chiang13_sim}, a structure with $\delta_g =12$ at $z\sim2.9$ has a 100\% probability of evolving into a galaxy cluster at $z=0$. \item[-] We estimated that this overdensity has a total mass of $M\sim8.1\times10^{14}M_{\odot}$ in a volume $13\times15\times17$ Mpc$^3$. According to \cite{chiang13_sim}, such an overdensity should collapse into a cluster of total mass $M_{z=0}\sim2.5 \times 10^{15} M_{\odot}$ at $z=0$. In the volume surveyed by VUDS at $2<z<3.5$ in the COSMOS field, we should have expected 0.12-0.3 proto-cluster of this kind. \item[-] The velocity dispersion of the 12 members is $\sigma_{los}=270 \pm 80 \kms$. We used light cones extracted from the Millennium Simulation to verify that this is lower (but consistent within $2\sigma$) than the typical velocity dispersion of the galaxies belonging to the same kind of proto-clusters at this redshift. This low value is consistent with the increase in $\sigma_{los}$ as time goes by \citep{eke98}. \item[-] In the light cones that we examined, the typical span in redshift of the galaxies belonging to proto-clusters at $z\sim2.9$, which will collapse into massive clusters at $z=0$, is $\Delta z \sim 0.02$. This value is much lower than the redshift bin often used to search for proto-clusters at this redshift. \item[-] The stacked spectra of the galaxies in the background of the overdensity show a significant absorption feature at the observed wavelength corresponding to the Ly$\alpha$ at the redshift of the structure ($\lambda=4736$\AA). We find that this absorption feature has a rest frame EW of $10.8\pm3.7$\AA, with a detection S/N of $\sim4$, when stacking only background galaxies without intervening sources at $z\sim2.9$ along their line of sight. We verified that this measurement is likely not to be due to noise fluctuations. Considering also the lower (but consistent) EW found using different samples of background galaxies (see Table \ref{EW_tab}), such an EW range corresponds to a column density N(HI) of the order of $3-20\times10^{19}$cm$^{-2}$. \item[-] We analysed the properties of the galaxies within the overdensity, and we compared them with a control sample at approximately the same redshift outside the overdensity. We could not find any statistically significant difference between the properties (stellar mass, SFR, sSFR, $NUV-r$, $r-K$) of the galaxies inside and outside the overdensity, but this result might be due to the lack of statistics, or possibly to the specific galaxy population sampled by VUDS, which could be less affected by environment than other populations not probed by the survey. \end{itemize} Simulations \citep{chiang13_sim} indicate that such an overdensity at $z\sim2.9$ is indeed a proto-cluster that, given the measured galaxy overdensity, will collapse to a (massive) cluster at $z=0$. For this reason, the detailed analysis of this proto-cluster represents a fundamental step in the comprehension of galaxy formation and evolution. For the properties of the galaxies within the proto-cluster, we plan to do a more detailed study when the VUDS selection function is assessed. A well-defined selection function will allow us to robustly quantify the average properties of the galaxies in such a dense environment and to compare them with the galaxies in the field at the same $z$ and with other overdensities found in the literature (e.g. \citealp{lemaux14}). The synergy of spectroscopy and multi-band photometry in next-generation surveys like Euclid will allow several proto-cluster structures to be identified thanks to combination of depth and large surveyed areas. In this respect, current surveys such as VUDS, are essential for characterising the properties of these structures and using them to predict their observability and optimal detection with future surveys. On the side of the detection of cold gas, the EW of the absorption line corresponding to the Ly$\alpha$ at $z\sim2.9$ implies a high column density (N(HI)$\sim10^{20}$cm$^{-2}$). This N(HI) value would be compatible with the scenario where the absorption is due to intervening cold streams of gas, which are falling (and feeding) into the halo's potential well of the galaxies in the proto-cluster \citep{fumagalli11,goerdt12}. In contrast, the stacked spectrum of the galaxies in the proto-cluster background also shows an absorption line corresponding to Si IV $\lambda\lambda1393,1402$ (blended) doublet at the redshift of the proto-cluster, detection that does not agree with the prediction that the cold flows are metal poor. Surely, the scenario of gas accretion by cold gas stream needs to be more robustly assessed from additional observational evidence. An exciting prospect is to look for the Ly$\alpha$ emission produced by gravitational energy released by cold gas flowing into the potential wells of galaxies (see e.g. \citealp{goerdt10}). | 14 | 3 | 1403.3691 |
1403 | 1403.0692_arXiv.txt | Line-of-sight magnetograms from the Helioseismic and Magnetic Imager (HMI) of the {\em Solar Dynamics Observatory} (SDO) are analyzed using a diagnostic known as the ``Magnetic Range of Influence,'' or MRoI. The MRoI is a measure of the length over which a photospheric magnetogram is balanced and so its application gives the user a sense of the connective length scales in the outer solar atmosphere. The MRoI maps and histograms inferred from the \sdo/HMI magnetograms primarily exhibit four scales: a scale of a few megameters that can be associated with granulation, a scale of a few tens of megameters that can be associated with super-granulation, a scale of many hundreds to thousands of megameters that can be associated with coronal holes and active regions, and a hitherto unnoticed scale that ranges from 100 to 250 megameters. We infer that this final scale is an imprint of the (rotationally-driven) giant convective scale on photospheric magnetism. This scale appears in MRoI maps as well-defined, spatially distributed, concentrations that we have dubbed ``g-nodes.'' Furthermore, using coronal observations from the Atmospheric Imaging Assembly (AIA) on \sdo{}, we see that the vicinity of these g-nodes appears to be a preferred location for the formation of extreme ultraviolet (EUV, and likely X-Ray) brightpoints. These observations and straightforward diagnostics offer the potential of a near-real-time mapping of the Sun's largest convective scale, a scale that possibly reaches to the very bottom of the convective zone. | The region of the solar atmosphere that is hidden from direct observation by its profound optical depth \-- the Sun's convective interior \-- is an ocean of boiling, bubbling plasma sustained from beneath by a rotating, nuclear furnace. The convective layer that forms the last third of the solar interior masks the process, or processes, which govern the production and perpetual eruption of the Sun's ubiquitous magnetic field \-- the magnetism which shapes the heliosphere, moderates the energy to fill it in addition to providing the energy necessary for life on our planet. Probing that ocean and understanding the magnetism of the solar interior, its production, evolution, and eventual destruction is a huge challenge and is the primary scientific focus of our community. The forced (magneto-)convection of the solar interior leaves visible tracers on the optical surface, the granular and super-granular scales that are a few and a few tens of megameters in diameter, respectively and the sunspots and active regions that can grow to several hundred megameters in size. However, granulation and supergranulation are only the tip of the proverbial (magnetic) iceberg \citep[e.g.,][]{Nordlund2009}. A tertiary scale of convection has tantalized the community, that of ``giant cell convection.'' A rotationally-forced convective scale \citep[e.g.,][]{Wilson1987} which was first realized in the pioneering calculations of \citet{1968ZA.....69..435S} and \citep{Gilman1975} that may play a critical role in the formation of sunspots and active regions \citep[e.g.,][]{Weber2012}. Giant cells are hypothesized to reach the bottom of the convective interior and span one to two hundred megameters in diameter \citep[see the review of][for an extensive discussion]{Miesch2005,Nordlund2009}, but they remained an elusive quarry to the observer with only scarce hints of their existence \citep[e.g.,][]{1998Natur.394..653B}. A recent analysis which tracked the motion of supergranules in long Dopplergram timeseries indicated that a large\--scale pattern could be discerned from the data and that the pattern demonstrated behavior consistent with that expected of giant convective cells \citep[][]{Hathaway2013}. It should be noted that contemporary investigations performed with high resolution observation and numerical simulations of solar surface convection strongly suggest that these apparently nested surface scales (with a strong emphasis towards the readily amenable granules and super-granules) reflect the multi-scale organization (or continuous spectrum) of motions present in the Sun's interior \citep[e.g.,][]{Berrilli2012, Orozco2012, Yelles2012, Yelles2014}. As such the third scale would indeed reflect a slower, longer-lived, convective scale present in the deepest interior \citep[e.g.,][]{Nordlund2009}. In the following sections we present the application of an analysis technique that was originally designed to investigate the range of connective length scales of the outer solar atmosphere that could be inferred from a simple analysis of photospheric magnetograms \-- the ``Magnetic Range of Influence'' \citep[MRoI;][]{2006ApJ...644L..87M}. The MRoI analysis demonstrates the presence of a tertiary convective scale and that is an ever-present in the magnetogram record of the Helioseismic and Magnetic Imager \citep[HMI;][]{2012SoPh..275..207S} of the {\em Solar Dynamics Observatory} (\sdo). Furthermore, we see that the magnetic elements which comprise the vertices of this larger scale convective pattern are a potential anchor-site for an ever-present feature of the solar corona---extreme-ultraviolet BrightPoints \citep[BPs;][]{1973SoPh...32...81V}. \begin{figure*} \epsscale{1.15} \plotone{f1} \caption{The full-disk \sdo{}/HMI line-of-sight magnetogram (A) and derived ``Magnetic Range of Influence'' (MRoI; B) map from May 16 2010. \label{f1}} \end{figure*} | We have seen that the quiescent photospheric magnetic field is composed of multiple connective scales. The observed scales range from a few megameters to those that are 100 \-- 250Mm in scale. We expect that the latter of these scales belongs to a spatially large, deep and hence slowly overturning convective flow \-- one that possibly reaches to the bottom of convection zone. Further, it would appear that photospheric line-of-sight magnetograms (and Dopplergrams) carry information about these nested scales in a non-trivial spatial mixture. It follows that the two cannot be easily disentangled without employing a technique like the MRoI. However, the ready visibility of a giant convective scale and its relatively straightforward identification could have a significant bearing on our ability to probe the variations of the deep solar interior and its long-term evolution. | 14 | 3 | 1403.0692 |
1403 | 1403.5886_arXiv.txt | {Galaxy clusters provide unique laboratories to study astrophysical processes on large scales and are important probes for cosmology. X-ray observations are currently the best means of detecting and characterizing galaxy clusters. Therefore X-ray surveys for galaxy clusters are one of the best ways to obtain a statistical census of the galaxy cluster population.} {In this paper we describe the construction of the REFLEX II galaxy cluster survey based on the southern part of the ROSAT All-Sky Survey. REFLEX II extends the REFLEX I survey by a factor of about two down to a flux limit of $1.8 \times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (0.1 - 2.4 keV).} {We describe the determination of the X-ray parameters, the process of X-ray source identification, and the construction of the survey selection function.} {The REFLEX II cluster sample comprises currently 915 objects. A standard selection function is derived for a lower source count limit of 20 photons in addition to the flux limit. The median redshift of the sample is $z = 0.102$. Internal consistency checks and the comparison to several other galaxy cluster surveys imply that REFLEX II is better than 90\% complete with a contamination less than 10\%.} {With this publication we give a comprehensive statistical description of the REFLEX II survey and provide all the complementary information necessary for a proper modelling of the survey for astrophysical and cosmological applications.} | Galaxy clusters are important astrophysical laboratories and cosmological probes (e.g. Sarazin 1986, Voit 2005, Borgani 2006, Vikhlinin et al. 2009, Allen et al. 2011, B\"ohringer 2011). While the latter references are based on X-ray observations of galaxy clusters, a lot of recent progress has also been made by optical cluster surveys (e.g. Rozo et al. 2010) and millimeter wave surveys using the Sunyaev-Zel'dovich effect (Reichardt et al. 2012, Benson et al. 2012, Marriage et al. 2011, Sehgal et al. 2011, PLANCK-Collaboration 2011). However, the currently most detailed view on the structure and the properties of clusters comes from X-ray observations. An X-ray survey is also the best means of efficiently detecting galaxy clusters as gravitationally bound and well evolved objects. X-ray observations thus provide statistically well defined, approximately mass selected cluster samples, since (i) X-ray luminosity is tightly correlated to mass (e.g. Pratt et al. 2009), (ii) bright X-ray emission is only observed for evolved clusters with deep gravitational potentials, (iii) the X-ray emission is highly peaked and projection effects are minimized, and (iv) for all these reasons the survey selection function can be accurately modeled. The X-ray emission originates in the several 10 Million degree plasma trapped in the cluster's gravitational potential well. In hydrostatic equilibrium, which is well approximated in most clusters that are not just in a stage of collision and merging, the intracluster density is tracing the equipotential surfaces. The observed X-ray luminosity is proportional to the square of the plasma density with usually a very weak temperature dependence in the X-ray energy band used by X-ray telescopes. The X-ray image provides, therefore, very important information on the mass distribution in the cluster and on its structure, even though we can see only one projection of the volume X-ray emission of the cluster (e.g. B\"ohringer et al. 2010). Systematic searches for galaxy clusters in X-rays are thus the currently most established prerequisite for comprehensive astrophysical studies, as well as for cosmological model testing. The {\sf ROSAT} All-Sky Survey (RASS, Tr\"umper 1993), is the only existing full sky survey conducted with an imaging X-ray telescope, providing a sky atlas in which one can search systematically for clusters in the nearby Universe. So far the largest, high-quality sample of X-ray selected galaxy clusters is provided by the {\sf REFLEX} Cluster Survey (B\"ohringer et al. 2001, 2004) based on the southern extragalatic sky of RASS at declination $\le 2.5$ degree. The quality of the sample has been demonstrated by showing that it can provide reliable measures of the large-scale structure (Collins et al. 2000, Schuecker et al. 2001a, Kerscher et al. 2001), yielding cosmological parameters (Schuecker et al. 2003a, b; B\"ohringer 2011) in good agreement within the measurement uncertainties with the subsequently published WMAP results (Komatsu et al. 2011). The {\sf REFLEX} data have also been used to study the galaxy velocity dispersion - X-ray luminosity relation (Ortiz-Gil et al., 2004), the statistics of Minkowski functionals in the cluster distribution (Kerscher et al. 2001), and to select statistically well defined subsamples like the HIFLUGCS (Reiprich \& B\"ohringer 2002) and REXCESS (B\"ohringer et al. 2007). The latter is particularly important as a representative sample of X-ray surveys to establish X-ray scaling relations (Croston et al. 2008, Pratt et al. 2009, 2010, Arnaud et al. 2010) and the statistics of the morphological distribution of galaxy clusters in X-rays (B\"ohringer et al. 2010). We also constructed a catalog of superclusters from the latest version of the REFLEX catalog comprising 164 superclusters including close pairs of clusters (Chon et al. 2013). Other galaxy cluster surveys based on the RASS comprise XBACs (X-ray-brightest Abell-type clusters; Ebeling et al. 1996), BCS ({\sf ROSAT} Brightest Cluster Sample) and eBCS survey (Ebeling et al. 1997,1998), RBS (RASS1 Bright Sample; De Grandi et al. 1999), NORAS (Northern {\sf ROSAT} ALL-SKY Survey Cluster Sample; B\"ohringer et al. 2000), the SGP survey (South Galactic Pole cluster survey; Cruddace et al. 2002), MACS (Most Massive Galaxy Clusters; Ebeling et al. (2000) and the CIZA survey (Clusters in the Zone of Avoidance; Ebeling et al. 2002, Kocevski et al. 2007). The RBS and SGP surveys were part of the early efforts to create the {\sf REFLEX} survey. In this paper we describe the extension of the {\sf REFLEX} survey from the previous flux limit of $3 \times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ in the 0.1 - 2.4 keV band to $1.8 \times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$. The number of clusters increases from 447 to 915 with this extension. We have already used this cluster sample to assess the power spectrum of the galaxy cluster distribution (Balaguera-Antolinez et al. 2010) with the interesting finding that the bias behavior of clusters in two-point statistics is exactly what is predicted by the theoretical statistical models. While this study was conducted when about 5\% of the galaxy cluster redshifts were still missing, we have now almost completed the spectroscopic follow-up observations in two observing campaigns in 2010 and 2011 (Chon \& B\"ohringer 2012) which leaves only 7 galaxy clusters without redshift information. This small number of missing redshifts will not significantly affect the statistics of the survey described in the present paper and in some first cosmological applications. Our plan for the publication of the full cluster catalog in the near future will be based on the completion of the redshift measurements. In the following we will use the term {\sf REFLEX II} for the extended cluster survey and {\sf REFLEX I} for the previous cluster survey. The paper is organized as follows. In chapter 2 we provide an overview on the global properties of the survey. Section 3 describes the determination of the X-ray parameters of the clusters and section 4 provides an overview on the source identification process, the selection of cluster candidates and the description of the follow-up observations. In section 5 we describe the construction of the survey selection function and in section 6 we derive various statistical properties of the {\sf REFLEX II} survey. In section 7 we compare our cluster detections to several other cluster surveys. Section 8 provides a discussion of the results and of the completeness and contamination of the {\sf REFLEX II} cluster sample. Section 9 comprises the summary and conclusions. For the derivation of distance dependent parameters we use a geometrically flat $\Lambda$-cosmological model with $\Omega_m = 0.3$ and $h_{70} = H_0/70$ km s$^{-1}$ Mpc$^{-1}$ = 1. All uncertainties without further specifications refer to 1$\sigma$ confidence limits. | The {\sf REFLEX II} sample with 915 galaxy clusters together with the northern {\sf NORAS II} sample of RASS detected clusters constitutes the largest statistically well defined sample of X-ray luminous galaxy clusters to date and will probably remain so until the exploitation of the {\sf eROSITA} All-Sky X-ray Survey (Predehl et al. 2011) becomes effective. The sample will therefore be important for astrophysical as well as cosmological studies. For these investigations the statistical properties of the survey and the survey selection process has to be well known. We have therefore made a large effort to construct a detailed three-dimensional survey selection function providing the luminosity selection limit as a function of the sky position and redshift. This provides a basis not only for the construction of simple zero order distribution functions like the X-ray luminosity function but also that of higher order functions like the N-point correlation functions, the power spectrum or e.g. Minkowski functionals. The survey selection function is not the only ingredient for the modeling of the survey in cosmological applications. A precise characterization of the measurement errors is equally important. Here it is the uncertainty of the measured photon counts which is practically proportional to the error of the measurement of the source flux and X-ray luminosity. As an improvement over {\sf REFLEX I} where we have used a mean flux error for the modeling, we have made an effort to derive a more detailed description of the flux uncertainty. Through a careful study of the dependence of the flux error on various cluster parameters, we have found that the most important parameter dependence of the flux error can be modelled as a function of $flux \times exposure$ as described in section 3. We have shown in section 8 that the best estimate for the completeness of the {\sf REFLEX II} catalog is of the order of 95\% while we expect a contamination level of about 5\%. Therefore it is not surprising that with this good quality of the catalog we can obtain a good, uncontaminated measure of the power spectrum of the spatial distribution of the clusters (Balaguera-Antolinez et al. 2010). To estimate how many clusters are still waiting to be detected in the {\sf REFLEX} area of the RASS among the more than 100 000 X-ray sources in total, we can use the same type of calculation as used for the photon number distribution shown in Fig.~\ref{fig17}. We estimate the total number of clusters expected to have more than six counts without any flux limit. With the low X-ray background of the RASS, a detection of six photons has still a high probability to be real. The result for the total number of these clusters in the {\sf REFLEX} region is about 9100. Thus we have only detected a small fraction of all clusters in the RASS so far. We have, however, reached some quality limit. Pursuing cluster detection to lower fluxes would increase the mean flux error to values higher than 20\% and also the cluster identification is getting notecibly more difficult. | 14 | 3 | 1403.5886 |
1403 | 1403.3849_arXiv.txt | {In more than four years of observation the Large Area Telescope on board the \emph{Fermi} satellite has identified pulsed $\gamma$-ray emission from more than 80 young or middle-aged pulsars, in most cases providing light curves with high statistics. Fitting the observed profiles with geometrical models can provide estimates of the magnetic obliquity $\alpha$ and of the line of sight angle $\zeta$, yielding estimates of the radiation beaming factor and radiated luminosity. Using different $\gamma$-ray emission geometries (Polar Cap, Slot Gap, Outer Gap, One Pole Caustic) and core plus cone geometries for the radio emission, we fit $\gamma$-ray light curves for 76 young or middle-aged pulsars and we jointly fit their $\gamma$-ray plus radio light curves when possible. We find that a joint radio plus $\gamma$-ray fit strategy is important to obtain $(\alpha,\zeta)$ estimates that can explain simultaneously detectable radio and $\gamma$-ray emission: when the radio emission is available, the inclusion of the radio light curve in the fit leads to important changes in the $(\alpha,\zeta)$ solutions. The most pronounced changes are observed for Outer Gap and One Pole Caustic models for which the $\gamma$-ray only fit leads to underestimated $\alpha$ or $\zeta$ when the solution is found to the left or to the right of the main $\alpha$-$\zeta$ plane diagonal respectively. The intermediate-to-high altitude magnetosphere models, Slot Gap, Outer Gap, and One pole Caustic, are favoured in explaining the observations. We find no apparent evolution of $\alpha$ on a time scale of $10^6$ years. For all emission geometries our derived $\gamma$-ray beaming factors are generally less than one and do not significantly evolve with the spin-down power. A more pronounced beaming factor vs. spin-down power correlation is observed for Slot Gap model and radio-quiet pulsars and for the Outer Gap model and radio-loud pulsars. The beaming factor distributions exhibit a large dispersion that is less pronounced for the Slot Gap case and that decreases from radio-quiet to radio-loud solutions. For all models, the correlation between $\gamma$-ray luminosity and spin-down power is consistent with a square root dependence. The $\gamma$-ray luminosities obtained by using the beaming factors estimated in the framework of each model do not exceed the spin-down power. This suggests that assuming a beaming factor of one for all objects, as done in other studies, likely overestimates the real values. The data show a relation between the pulsar spectral characteristics and the width of the accelerator gap. The relation obtained in the case of the Slot Gap model is consistent with the theoretical prediction.} \authorrunning{Pierbattista et al. 2013} \titlerunning{Magnetic obliquity and line of sight constraints} | The advent of the Large Area Telescope \citep[LAT,][]{aaa+09a} on the \emph{Fermi} satellite has significantly increased our understanding of the high-energy emission from pulsars. After more than four years of observations the LAT has detected pulsed emission from more than 80 young or middle-aged pulsars, collecting an unprecedented amount of data for these sources \citep[][]{2PC}. This has allowed the study of the collective properties of the $\gamma$-ray pulsar population \citep{pie10,wr11,twc11,pghg12} and of the pulse profiles. The light-curve analysis can be approached by studying the number of peaks and morphology or by modelling the $\gamma$-ray profiles to estimate pulsar orientations and constrain the model that best describes the observations. The first type of analysis has been performed by \cite{wrwj09} and \cite{pie10}, who studied light-curve peak separation and multiplicities in light of intermediate and high-altitude gap magnetosphere models. The second type of analysis has been performed for a small set of pulsars by \cite{rw10} and \cite{pie10} for young and middle-aged pulsars, and \cite{vhg09} for millisecond pulsars. They used the simulated emission patterns of proposed models to fit the observed light curves and estimate the magnetic obliquity angle $\alpha$ (the angle between the pulsar rotational and magnetic axes) and the observer line of sight angle $\zeta$ (the angle between the observer direction and the pulsar rotational axis), showing that the outer magnetosphere models are favoured in explaining the pulsar light curves observed by \emph{Fermi}. What these first studies suggest is that with the new high-statistics of the LAT pulsar light curves, fitting the observed profiles with different emission models has become a powerful tool to give estimates of the pulsar orientation, beaming factor, and luminosity, and to constrain the geometric emission models. After discovery of the pulsed high-energy emission from the Crab pulsar \citep{mbc+73}, emission gap models were the preferred physical descriptions of magnetospheric processes that produce $\gamma$-rays. These models predict the existence of regions in the magnetosphere where the Goldreich \& Julian force-free condition \citep{gj69} is locally violated and particles can be accelerated up to a few TeV. Three gap regions were identified in the pulsar magnetosphere: the Polar Cap region \citep{stu71}, above the pulsar polar cap; the Slot Gap region \citep{aro83b}, along the last closed magnetic field line; the Outer Gap region \citep{chr86}, between the null charge surface and the light cylinder. \cite{dhr04} calculated the pulsar emission patterns of each model, according to the pulsar magnetic field, spin period, $\alpha$, and gap width and position. The \cite{dhr04} model is based on the assumptions that the magnetic field of a pulsar is a vacuum dipole swept-back by the pulsar rotation \citep{deu55} and that the $\gamma$-ray emission is tangent to the magnetic field lines and radiated in the direction of the accelerated electron velocity in the co-rotating frame. The emission pattern of a pulsar is then obtained by computing the direction of $\gamma$-rays from a gap region located at the altitude range characteristic of that model. Note that the number of radiated $\gamma$-rays depends only on the emission gap width and maximum emission radius, which are assumed parameters. The aim of this paper is to compare the light curves of the young and middle-aged LAT pulsars listed in the second pulsar catalog \citep[][hereafter PSRCAT2]{2PC} with the emission patterns predicted by theoretical models. We use the \cite{dhr04} geometric model to calculate the radio emission patterns according to radio core plus cone models \citep{gvh04,sgh07,hgg07,pghg12}, and the $\gamma$-ray emission patterns according to the Polar Cap model \citep[PC,][]{mh03}, the Slot Gap model \citep[SG,][]{mh04a}, the Outer Gap model \citep[OG,][]{crz00}, and an alternative formulation of the OG model that differs just in the emission gap width and luminosity formulations, the One Pole Caustic \citep[OPC,][]{rw10,wrwj09} model. We use them to fit the observed light curves and obtain estimates of $\alpha$, $\zeta$, outer gap width $w_\mathrm{OG/OPC}$, and slot gap width $w_\mathrm{SG}$, as well as the ensuing beaming factor and luminosity. Using these estimates, we study the collective properties of some non-directly observable characteristics of the LAT pulsars, namely their beaming factors, $\gamma$-ray luminosity, magnetic alignment, and correlation between the width of the accelerator gap and the observed spectral characteristics. For each pulsar of the sample and each model, the estimates of $\alpha$ and $\zeta$ we obtain represent the best-fit solution in the framework of that specific model. We define the \emph{optimum-solution} as that solution characterised by the highest log-likelihood value among the four emission models, and we define the \emph{optimum-model} as the corresponding model. Hereafter we will stick to this nomenclature in the descriptions of the fit techniques and in the discussion of the results. The radio and/or $\gamma$-ray nature of the pulsars of our sample have been classified according to the flux criterion adopted in PSRCAT2: radio-quiet (RQ) pulsars, with radio flux detected at 1400 MHz $S_{1400}<30\mu$Jy and radio-loud (RL) pulsars with $S_{1400}>30\mu$Jy. The $30\mu$Jy flux threshold was introduced in PSRCAT2 to favour observational characteristics instead of discovery history in order to have more homogeneous pulsar samples. Yet, radio light curves were available for 2 RQ pulsars, J0106$+$4855 and J1907$+$0602, that show a radio flux $S_{1400}<30\mu$Jy (PSRCAT2). We include these two radio-faint (RF) pulsars in the RQ sample and the results of their joint $\gamma$-ray plus radio analysis are given in Appendix \ref{JointFits_RQ2RL}. The outline of this paper is as follows. In Section \ref{Data} we describe the data selection criteria adopted to build the $\gamma$-ray and radio light curves. In Section \ref{Simulation} we describe the method we use to calculate the pulsed emission patterns and light curves. Sections \ref{Individual gamma-ray fit} and \ref{Fitting both the gamma-ray and radio emission} describe the fitting techniques used for the RQ and RL pulsars, respectively. The results are discussed in Section \ref{Results}. In Appendix \ref{GoodFitMethod} we describe the method used to give an estimate of the relative goodness of the fit solutions. In Appendix \ref{PopDis} we show further results obtained from the pulsar population synthesis study of \cite{pghg12} that we will compare with results obtained in Sections \ref{A-Z best solutions plane} and \ref{HighECut}. Appendices \ref{GammaFitRes}, \ref{JointFitRes}, and \ref{JointFits_RQ2RL} show, for each model, the best-fit $\gamma$-ray light curves for RQ LAT pulsars, the best-fit $\gamma$-ray and radio light curves for RL LAT pulsars, and the best-fit $\gamma$-ray and radio light curves of two RQ-classified LAT pulsars for which a radio light curve exists. | We have selected a sample of young and middle-aged pulsars observed by the LAT during three years and described in PSRCAT2. We have fitted their $\gamma$-ray and radio light curves with simulated $\gamma$-ray and radio emission patterns. We have computed the radio emission beam according to \cite{sgh07} and we have used the geometrical model of \cite{dhr04} to simulate the $\gamma$-ray emission according to four gap models, PC, \citep{mh03}, SG, \citep{mh04a}, OG, \citep{crz00} and OPC \citep{rw10,wrwj09}. Each emission pattern has been described by a series of phase-plots, evaluated for the pulsar period, magnetic field, and gap width, and for the whole $\alpha$ interval sampled every degree. These phase-plots predict the pulsar light curve as a function of $\zeta$. The simulated phase-plots have been used to fit the observed radio and $\gamma$-ray light curves according to two different schemes: a single fit to the $\gamma$-ray profiles of RF and RQ objects and a joint fit to the $\gamma$-ray and radio light curves of RL pulsars. The individual fit to the $\gamma$-ray profiles has been implemented using a $\chi^2$ estimator and light curves binned both in FCBin and RBin. The comparison of the results obtained with the two methods shows that the $\chi^2$ fit with FCBin light curves yields the closest match between the observations and modelled profiles. We use the latter to give $\alpha$ and $\zeta$ estimates for the RQ and RF LAT pulsars and we use the RBin fit to evaluate the systematic uncertainties induced by the fitting method. The joint $\gamma$-ray plus radio fit of RL pulsars uses RBin radio light curves and FCBin $\gamma$-ray light curves with a $\chi^2$ estimator. The log-likelihood maps in $\alpha$ and $\zeta$ obtained from the radio-only and $\gamma$-ray-only fits were summed to produce the joint solution. Two options were considered to couple the high signal-to-noise ratio of the radio data to the much lower signal-to-noise ratio of the $\gamma$-ray profiles and the solution characterised by the highest log-likelihood value was selected. The systematic errors on $(\alpha,\zeta)$ for the RL pulsars have been obtained by studying the difference between the solutions obtained with the two joint fit coupling schemes. We have obtained new constraints on $\alpha$ and $\zeta$ for 33 RQ, 2 RF, and 41 RL $\gamma$-ray pulsars. We have studied how the $(\alpha,\zeta)$ solutions of RL pulsars obtained by fitting only the $\gamma$-ray light curves change by including the radio emission in the fit. We have used the $\alpha$ and $\zeta$ solutions to estimate several important pulsar parameters: gap width, beaming factor, and luminosity. We have also investigated some relations between observable characteristics and intrinsic pulsar parameters, such as $\alpha$ as a function of age and the spectral energy cut-off and index in $\gamma$-rays as a function of the gap width. We find no evidence for an evolution of the magnetic obliquity over the $\sim10^6$ yr of age span in the sample, but we find an interesting apparent change in the $\gamma$-ray spectral index $\Gamma$ and high-energy cutoff $E_\mathrm{cut}$ associated with changes in the gap widths. We have found that a multi-wavelength fit of $\gamma$-ray and radio light curves is important in giving a pulsar orientation estimate that can explain both radio and $\gamma$-ray emission. The PC emission geometry explains only a small fraction of the observed profiles, in particular for the RL pulsars, while the intermediate to high SG and OG/OPC models are favoured in explaining the pulsar emission pattern of both RQ and RL LAT pulsars. The fact that none of the assumed emission geometries is able to explain all the observed LAT light curves suggests that the true $\gamma$-ray emission geometry may be a combination of SG and OG and that we detect the respective light curves for different observer viewing angles. Comparison of the $\alpha$ and $\zeta$ solutions obtained by fitting only the $\gamma$-ray profiles of RL pulsars and both their $\gamma$-ray and radio profiles suggests that in the OG and OPC models, $\alpha$ or $\zeta$ are underestimated when one does not account for radio emission. When the $\gamma$-only solution is to the right of the radio diagonal in the $\alpha$-$\zeta$ plane, $\zeta$ migrates toward higher values while $\alpha$ keeps quite stable and \emph{vice versa} when the $\gamma$-only solution is to the left of the radio diagonal. The beaming factors found for the RQ and RL objects are consistent with the distributions obtained in the population study of \cite{pghg12}. For all the models we observe a large scatter of the beaming factors with $\dot{E}$, which is reduced for RL pulsars compared to RQ pulsars, except for the SG. This is because RQ pulsars are viewed at lower $\alpha$ and $\zeta$, and OG and OPC beams shrink towards the spin equator with decreasing $\dot{E}$ while SG beams do not. The low $f_{\Omega}$ values found for the PC reflect the narrow geometry of the PC beams. The $f_{\Omega}$ values for the SG appear to be fairly stable around 1 over 4 decades in $\dot{E}$. We find also little evolution for the OG and OPC beaming factors of RQ objects which gather around 0.25 and 0.39, respectively. Larger averages are obtained for the RL objects (0,68 for OG and 0,86 for OPC) with no evolution with $\dot{E}$ for the OPC case and some hint of an increase with $\dot{E}$ in the OG case. The fact that the majority of the pulsars exhibit an $f_{\Omega}$ estimate less than unity in all models suggests that the isotropic luminosities ($f_{\Omega}$ = 1) often quoted in other studies are likely to overestimate the real values. For all the models a power law relation consistent with $L_{\gamma}~\appropto~\dot{E}^{0.5}$ is observed for both RQ and RL pulsars. In contrast with PSRCAT2 we do not obtain any $\gamma$-ray luminosities significantly higher than $\dot{E}$. Since the only difference between the luminosity computation here and that of PSRCAT2 is in the $f_{\Omega}$ value (assumed equal to one in the catalog), the excessively high luminosities obtained in the catalog probably result from a too high beaming factor. We have studied the consistency of the geometric $\gamma$-ray luminosity, $L_\mathrm{geo}$, obtained in this paper and the $\gamma$-ray luminosity computed in the framework of radiative gap-models, $L_\mathrm{rad}$. We found that $L_\mathrm{geo}$ overestimate $L_\mathrm{rad}$ of 2-3 order of magnitude for the RQ and RL SG pulsar and for RQ OG pulsars while the $L_\mathrm{geo}$ of RL OG objects are more consistent with their $L_\mathrm{rad}$ values while showing higher dispersion in $L_\mathrm{geo}$. For both RQ and RL OPC objects, $L_\mathrm{rad}$ is consistent with the $L_\mathrm{rad}$ estimates. These OG and SG geometric-radiative luminosity disagreements are due to inconsistencies in the formulation of the geometrical and radiative aspects of the $\gamma$-ray pulsar emission, rise the problem of formulating geometrical models more based on the actual pulsar electrodynamics in the framework of each gap model, and points to fundamental shortcomings of these electrodynamic gap models. We find a correlation between $E_\mathrm{cut}$ and $\Gamma$ of the $\gamma$-rays and the accelerator gap width in the magnetosphere. The relation is consistent with the SG prediction $E_\mathrm{cut}\appropto w_\mathrm{SG}^{-0.25}$ just for the RL objects while the more approximated predictions formulated for OG and OPC models are not consistent with the observations. This $E_\mathrm{cut}$ and $\Gamma$ versus gap width proportionality is important because it connects the observed spectral information and the non observable size of the gap region on the basis of the light-curve morphology alone. | 14 | 3 | 1403.3849 |
1403 | 1403.3078_arXiv.txt | We report on a \nustar\ and \xmm\ program that has observed a sample of three extremely luminous, heavily obscured \wise-selected AGN at $z \sim 2$ in a broad X-ray band ($0.1 - 79$~keV). The parent sample, selected to be faint or undetected in the \wise\ $3.4 \mu$m ($W1$) and $4.6 \mu$m ($W2$) bands but bright at $12 \mu$m ($W3$) and $22 \mu$m ($W4$), are extremely rare, with only $\sim 1000$ so-called ``$W1W2$-dropouts'' across the extragalactic sky. Optical spectroscopy reveals typical redshifts of $z \sim 2$ for this population, implying rest-frame mid-IR luminosities of $\nu L_\nu(6 \mu{\rm m}) \sim 6 \times 10^{46}\, {\rm erg}\, {\rm s}^{-1}$ and bolometric luminosities that can exceed $L_{\rm bol} \sim 10^{14}\, L_\odot$. The corresponding intrinsic, unobscured hard X-ray luminosities are $L(2-10\, {\rm keV}) \sim 4 \times 10^{45}\, {\rm erg}\, {\rm s}^{-1}$ for typical quasar templates. These are amongst the most luminous AGN known, though the optical spectra rarely show evidence of a broad-line region and the selection criteria imply heavy obscuration even at rest-frame $1.5\, \mu{\rm m}$. We designed our X-ray observations to obtain robust detections for gas column densities $N_{\rm H} \leq 10^{24}\, {\rm cm}^{-2}$. In fact, the sources prove to be fainter than these predictions. Two of the sources were observed by both \nustar\ and \xmm, with neither being detected by \nustar\ ($f_{\rm 3-24\ keV} \simlt 10^{-13}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$), and one being faintly detected by \xmm\ ($f_{\rm 0.5-10\ keV} \sim 5 \times 10^{-15}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$). A third source was observed only with \xmm, yielding a faint detection ($f_{\rm 0.5-10\ keV} \sim 7 \times 10^{-15}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$). The X-ray data require gas column densities $N_{\rm H} \simgt 10^{24}\, {\rm cm}^{-2}$, implying the sources are extremely obscured, consistent with Compton-thick, luminous quasars. The discovery of a significant population of heavily obscured, extremely luminous AGN does not conform to the standard paradigm of a receding torus, in which more luminous quasars are less likely to be obscured. If a larger sample conforms with this finding, then this suggests an additional source of obscuration for these extreme sources. | The {\it Wide-field Infrared Survey Explorer} (\wise) \citep{Wright:10} is an extremely capable and efficient black hole finder. As demonstrated in selected fields by \spitzer\ \citep[\eg,][]{Stern:05, Donley:12}, the same material that obscures AGN at UV, optical and soft X-ray energies is heated by the AGN and emits strongly at mid-IR wavelengths. The all-sky \wise\, survey identifies millions of obscured and unobscured quasars across the full sky \citep[\eg,][]{Stern:12, Assef:13}, as well as very rare populations of extremely luminous, heavily obscured AGN. In terms of the latter, the \wise\, extragalactic team has been pursuing sources that are faint or undetected in \wise\, $W1$ (3.4~$\mu$m) and $W2$ (4.6~$\mu$m), but are bright in $W3$ (12~$\mu$m) and $W4$ (22~$\mu$m). We refer to this population as $W1W2$-dropouts \citep{Eisenhardt:12}. This is a very rare population; selecting to a depth of 1~mJy at 12~$\mu$m, there are only $\sim 1000$ such sources across the extragalactic sky ($\sim 1$ per $30\ {\rm deg}^2$). These objects are undetected by {\it ROSAT} and tend to be optically faint ($r \simgt 23$), below the detection threshold of SDSS. We have obtained spectroscopic redshifts for $> 100$ $W1W2$-dropouts thus far, consistently finding redshifts $z \simgt 2$, with our current highest redshift source at $z = 4.6$ \citep{Eisenhardt:14}. Approximately half of the sources show clear type-2 AGN signatures in the optical spectra, with the other half typically showing only Ly$\alpha$ emission, sometimes extended, which could be due to star formation and/or AGN activity \citep{Bridge:13}. The lack of a far-IR peak in their broad-band SEDs suggests the dominant energy input for this population comes from a heavily obscured AGN and not extreme starbursts \citep[\eg,][]{Eisenhardt:12, Wu:12}. Related high-luminosity sources selected from the \wise\ satellite have also recently been reported by \citet{Weedman:12} and \citet{Alexandroff:13}, while several teams have identified less rare, less luminous sources from \spitzer\ surveys with less extreme colors \citep[\eg,][]{Dey:08, Fiore:09}. Here we report on the first targeted X-ray follow-up of $W1W2$-dropouts. We observed two sources with both the {\it Nuclear Spectroscopic Telescope Array} \citep[\nustar;][]{Harrison:13} and \xmm\ \citep{Jansen:01}; a third source was only observed by \xmm. Unless otherwise specified, we use Vega magnitudes throughout and adopt the concordance cosmology, $\Omega_M = 0.3$, $\Omega_\Lambda = 0.7$ and $H_0 = 70\, \kmsMpc$. | We designed our X-ray integration times to provide robust detections for (i) typical intrinsic AGN SEDs, and (ii) gas column densities $N_{\rm H} \simlt 10^{24}\, {\rm cm}^{-2}$. None of the three sources was strongly detected, implying that at least one of the assumptions in our experimental design does not hold. \nustar\ has now observed a range of obscured AGN, from famous, local sources such as Mrk~231 \citep{Teng:14}, Circinus \citep{Arevalo:14}, NGC~424 \citep{Balokovic:14}, and NGC~4945 \citep{Puccetti:14}, to higher redshift obscured quasars at $z \sim 0.5$ from SDSS \citep{Lansbury:14} and $z \sim 2$ in the ECDFS \citep{DelMoro:14}. A recurring theme of these observations is that several AGN which are extremely luminous at certain wavelengths, such as in the mid-IR or [\ion{O}{3}]~$\lambda 5007$, remain faint at X-ray energies. For some objects, this is the case even for the more penetrating hard X-rays $> 10$~keV. For some sources, such as optically bright ($B \simlt 16$) broad-absorption line (BAL) quasars from the Palomar-Green (PG) survey \citep{Schmidt:83} and Mrk~231, we consider, and sometimes even favor attributing the hard X-ray faintness to intrinsic X-ray weakness \citep{Luo:13, Teng:14}. Such intrinsic X-ray weakness seems the most plausible scenario when the AGN appears unobscured at certain wavelengths, through strong UV continuum emission, broad emission lines, and/or weak X-ray spectra that are well described by a moderately absorbed power-law. For a more detailed description of the intrinsically X-ray weak scenario, see \citet{Luo:13} and \citet{Teng:14}. For the three \wise-selected AGN discussed here, we instead favor interpreting the X-ray faintness as being due to a typical AGN seen through extremely high absorbing columns, consistent with their mid-IR SEDs and optical spectra. Indeed, the X-ray column constraints from Fig.~\ref{fig:LxLir} are broadly consistent with the mid-IR measurements given typical luminous AGN gas-to-dust ratios from \citet{Maiolino:01}. What is more surprising is that these sources, amongst the bolometrically most luminous AGN known, appear heavily obscured. Various observations have shown that more luminous AGN are less likely to be obscured \citep[\eg,][]{Ueda:03, Simpson:05, Assef:13}. This is consistent with the ``receding torus model'', first proposed by \citet{Lawrence:91}, in which the height of the torus is independent of luminosity while the the inner radius of the torus, corresponding to the distance at which dust reaches its sublimation temperature, increases with luminosity. Therefore, in this model, more luminous AGN have more sightlines into the nucleus and thus have a lower likelihood of being obscured. The $W1W2$-dropout population is a rare population, with a surface density of just one source per $\sim 30\, {\rm deg}^2$ in the extragalactic sky. The mid-IR luminosities imply intrinsic X-ray luminosities of a few $\times 10^{45}\, {\rm erg}\, {\rm s}^{-1}$ from the relations of \citet{Lutz:04} and \citet{Gandhi:09}. \citet{Just:07} report on X-ray follow-up of the most luminous quasars in the SDSS available at the time, finding 34 quasars across 4188 deg$^2$, or one source per $\sim 120\, {\rm deg}^2$. The X-ray luminosities of this luminous quasar sample prove comparable to the expected intrinsic X-ray luminosities from the \wise-selected sample, implying the surprising discovery of comparable numbers of obscured and unobscured quasars at the top of the luminosity function. \citet{Assef:14} and \citet{Tsai:14} present more detailed comparisons between the $W1W2$-dropout and luminous unobscured quasar populations. The discovery of a significant population of heavily obscured, extremely luminous AGN does not conform to the simple receding torus model, suggesting an additional source of obscuration. Indeed, the models of \citet{Draper:10} predict that Compton-thick AGN should be more common at higher redshift because of the high fueling rates of quasars require significant gas reservoirs. This is consistent with models showing that mergers may be more prominent in fueling AGNs at $z \sim 2$ \citep[\eg,][]{Hopkins:08, Draper:12}. Indeed, these models also predict that the unified model of AGN will break down for high-luminosity AGNs at $z \simgt 1$ because the obscuration is not confined to the nucleus \citep[see also][]{Draper:11}. Further investigations into this interesting population are clearly warranted. Deeper X-ray observations, achieving robust detections rather than faint and non-detections, would be valuable. For example, assuming the obscuration is from a nuclear torus, deeper X-ray observations should detect narrow, reflected Fe K$\alpha$ fluorescent emission at rest-frame 6.4~keV \citep[\eg,][]{Nandra:07}; a significant non-detection could point towards obscuration on larger scales than the torus. Such large-scale obscuration could also be probed by high-resolution imaging, such as far-IR observations with ALMA, and, eventually, mid-IR observations with {\it JWST}. Near-IR spectroscopy could also look for reddening in the AGN narrow-line region \citep[\eg,][]{Brand:07}, which is expected to be significantly larger than the torus. A clearer understanding of the geometry of the obscuring region combined with an improved reckoning of the $W1W2$-dropout space density will enable us to better place this population within the context of AGNs and galaxy evolution. | 14 | 3 | 1403.3078 |
1403 | 1403.6694_arXiv.txt | { Primordial or big bang nucleosynthesis (BBN) is one of the three historical strong evidences for the big bang model. The recent results by the {\it{Planck}} satellite mission have slightly changed the estimate of the baryonic density compared to the previous WMAP analysis. This article updates the BBN predictions for the light elements using the cosmological parameters determined by {\it{Planck}}, as well as an improvement of the nuclear network and new spectroscopic observations. There is a slight lowering of the primordial Li/H abundance, however, this lithium value still remains typically 3 times larger than its observed spectroscopic abundance in halo stars of the Galaxy. According to the importance of this "lithium problem", we trace the small changes in its BBN calculated abundance following updates of the baryonic density, neutron lifetime and networks. In addition, for the first time, we provide confidence limits for the production of \six, \neu, \onz\ and CNO, resulting from our extensive Monte Carlo calculation with our extended network. A specific focus is cast on CNO primordial production. Considering uncertainties on the nuclear rates around the CNO formation, we obtain $\rm{CNO/H} \approx (5-30)\times10^{-15}$. We further improve this estimate by analyzing correlations between yields and reaction rates and identified new influential reaction rates. These uncertain rates, if {\em simultaneously} varied could lead to a significant increase of CNO production: $\rm{CNO/H}\sim10^{-13}$. This result is important for the study of population III star formation during the dark ages. } \begin{document} | There are three historical observational evidences for the big bang model: the cosmic expansion, the Cosmic Microwave Background (CMB) radiation and primordial or big bang nucleosynthesis (BBN). Today, they are complemented by a large number of evidences in particular from the properties of the large scale structures (see e.g. \citet{pubook} for a textbook description). BBN predicts the primordial abundances of the ``light cosmological nuclei'': \hli\ that are produced during the first 20 min after the big bang when the Universe was dense and hot enough for nuclear reactions to take place (see e.g. \citet{Ste07,Ioc09,fields11} for recent reviews). The comparison of the calculated and observed abundances shows an overall good agreement except for the \sep. The essential cosmological parameter of the model is the baryonic density $\Omega_{\rm b}$. It is related to the baryon to photon ratio, $\eta\equiv n_{\rm b}/n_{\gamma}=2.738\times10^{-8}\;\Omega_{\mathrm{b}}h^2$ (see the appendix) that remains constant during the expansion after the electron--positron annihilation. \obh\ is now well measured from the angular power spectrum of the CMB temperature anisotropies. A precise value for this, previously free, parameter was provided by the Wilkinson Microwave Anisotropy Probe (WMAP9) satellite, $\Omega_{\rm b}h^2=0.02243\pm0.00055$,~ ("Nine-year (MASTER)", \citet{WMAP9}) while the recent {\it{Planck}} mission updated it to $\Omega_{\rm b}h^2$=0.02218$\pm$0.00026 ("{\it{Planck}}+lensing+WP+highL", \citet{Planck13}). This value is chosen because it includes all the last cosmological constraints. We calculate here the \hli\ primordial abundances by Monte Carlo, using our extended 424 nuclear reaction network \citep{Coc12a}, also taking into account the updated value of the neutron lifetime \cite{PDG12}. In \sbbn, only traces of other isotopes are produced: \six, \neu, \dix, \onz\ and CNO. The CNO abundance is of peculiar interest since it may affect Pop III stellar evolution in the first structures of the Universe. The value which could impact this evolution is estimated to be 10$^{-11}$ \cite{Cas93} or even as low as 10$^{-13}$ (in number of atoms relative to hydrogen, CNO/H) for the less massive stars \cite{Eks08}. In this context, it is important to evaluate carefully the BBN CNO abundance. In our previous work \citep{Coc12a} we obtained a much lower value CNO/H=$0.7\times10^{-15}$ but no upper nor lower limit (see also Ref. \cite{Ioc07}). In this paper, we use the results of our Monte Carlo calculations $i$) to estimate the uncertainties on the BBN production of the minor isotopes, and in particular of CNO and $ii$) analyze the correlations between reaction rates and isotopic abundances to identify potentially important reactions that were not identified in our previous sensitivity analysis. We show that by calculating correlations, we find important reactions that were overlooked in sensitivity studies changing one reaction at a time. This is crucial because the level of the CNO abundance plays a key role in the evolution of the first stars. | 14 | 3 | 1403.6694 |
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1403 | 1403.6377_arXiv.txt | The Pluto-Charon system, likely formed from an impact, has reached the endpoint of its tidal evolution. During its evolution into the dual-synchronous state, the equilibrium tidal figures of Pluto and Charon would have also evolved as angular momentum was transferred from Pluto's spin to Charon's orbit. The rate of tidal evolution is controlled by Pluto's interior physical and thermal state. We examine three interior models for Pluto: an undifferentiated rock/ice mixture, differentiated with ice above rock, and differentiated with an ocean. For the undifferentiated case without an ocean, the Pluto-Charon binary does not evolve to its current state unless its internal temperature $T_i>200$ K, which would likely lead to strong tidal heating, melting, and differentiation. Without an ocean, Pluto's interior temperature must be higher than 240 K for Charon to evolve on a time scale less than the age of the solar system. Further tidal heating would likely create an ocean. If \emph{New Horizons} finds evidence of ancient tidally-driven tectonic activity on either body, the most likely explanation is that Pluto had an internal ocean during Charon's orbital evolution. | \label{sec:intro_section} As the \emph{New Horizons} spacecraft travels toward the Pluto/Charon system, our knowledge about the system's physical properties \citep{Olkin2003, Gulbis2006, Tholen2008}, mode of formation \citep{Canup2005,Canup2011}, and surface composition (e.g., \citealt{Olkin2003, BrownCalvin2000, BuieGrundy2000, Cook2007}) continues to grow. The general properties of the Pluto and Charon system, including its dynamical state and the Pluto-to-Charon mass ratio, suggest that Charon may have formed due to a collision between two like-sized precursor objects \citep{Canup2005,Canup2011}. Careful examination of the system in preparation for the \emph{New Horizons} mission has led to the discovery of four additional small moons, Nix, Hydra \citep{Weaver2006}, Styx, and Kerberos \citep{Showalter2011,Showalter2012}. The moons are in mean motion resonances with Charon \citep{Weaver2006, Stern2006, Buie2006, Showalter2013}, which could imply that they formed by accretion of debris after the Charon-forming impact or were captured \citep{Ward2006, KenyonBromley2013}. The Pluto/Charon system is unique among major bodies in the solar system because it has reached the endpoint of its dynamical evolution: the so-called ``dual synchronous'' state in which Charon's orbital period, spin period, and Pluto's rotation period are equal; Charon currently orbits Pluto at $a_{c} \sim 16.4 R_P$, where $R_P=1147$ km is Pluto's radius \citep{Tholen2008}. As the orbit evolved to its present state and Pluto's spin rate changed to match Charon's orbital period, changes in the equilibrium tidal and rotational figures of the two bodies may have left their mark as systems of tectonic features on the surfaces of each body. Orbit and spin evolution of the Pluto/Charon system is driven by the raising and lowering of tidal bulges on each body. The tidal bulges exert torques which change the semi-major axes of the orbits and the spin rates. The mechanical energy associated with the periodic raising and lowering of the bulges is dissipated as heat in the bodies' interiors. Changes in bulge height and energy dissipation are thought to drive endogenic resurfacing and tectonic activity on many of the icy satellites of the outer solar system (see, e.g., \citealt{SchubertSatsBook, PealeSatsBook,EuropaJupBook} for discussion). The magnitude of stresses arising from tidal evolution depends on the interior structures of the bodies and the frequency of tidal flexing; the frequencies in turn depend on the Pluto/Charon distance and the spin periods. The main source of stress accumulation in the Pluto/Charon system would be figure changes associated with transfer of momentum from Pluto's spin to Charon's orbit. If the stresses exerted on the surfaces of the bodies exceed the nominal yield stress of ice, we consider that tectonic activity may have occurred. The presence or absence of tectonic features, along with their distribution and orientations, may provide clues about the early evolution of the system \citep{CollinsPappalardoLPSC2000}. Thus, images of the surfaces of Pluto and Charon from \emph{New Horizons} could yield clues about the post-impact interior state of Pluto and the initial orbital distance of Charon. In successful hydrodynamical simulations of the Charon-forming impact, Charon is launched into an eccentric orbit around a rapidly spinning Pluto (cf. \citealt{DobroPlutoBook}); its initial orbital semi-major axis after the impact, $a_o$, ranges from $\sim 3.7R_P$ to $21 R_P$ \citep{RobinPers}. Over a time scale $\tau_{synch} \sim 35 (Q_c/100)$ years \citep{DobroPlutoBook}, where $Q_c\sim 100$ is a nominal estimate of Charon's tidal quality factor, Charon evolves to a synchronous state in which its spin period and orbital period are equal (similar to the Earth-Moon system). The final migration from the synchronous state to the dual-synchronous state takes much longer, $\Delta t_{evol} \approx 200 (Q_p/100)(10^{-3}/k_{2,p})$Myr, where $Q_p$ is Pluto's tidal quality factor, and $k_{2,p}$ is the degree-2 Love number that describes how Pluto's gravitational potential changes in response to the tides raised on Pluto by Charon. In this work, we calculate the migration timescale and stresses generated in Pluto's lithosphere due to the orbital evolution of the system from its initial post-impact state to its present dual-synchronous state. We determine the Love numbers of Pluto as a function of its interior structure, temperature, and the time scale of deformation. The Love numbers are used to estimate its $Q$ immediately post-impact, and to determine the magnitude of tidal deformation in a rapidly spinning Pluto. The Love numbers are also used to constrain the magnitude of deformation and stresses built up in Pluto's lithosphere. The $Q$ values based on realistic interior structures for Pluto are used to constrain the orbital evolution time scale of the system and the rate of deformation associated with decreasing Pluto's tidal bulge as Charon recedes. Assuming nominal parameters describing the brittle and ductile behavior of water ice (e.g., Table \ref{table:lovenumber_params}), we determine the conditions under which the induced stresses and deformation rates can fracture Pluto's surface, and determine how that likelihood varies as a function of Pluto's thermal state post-impact. | \label{sec:conclusions} The \emph{New Horizons} mission has sparked new interest about the thermal evolution of Pluto and its moon Charon, in particular, the possibility of tectonic resurfacing on both bodies and the likelihood that these bodies may have had liquid water oceans, either now or in the past. In this work, we have used simple tidal models to calculate the magnitude of stresses expected on the surfaces of Pluto and Charon during their evolution into the dual-synchronous state. During this time, Pluto de-spins as Charon recedes from Pluto, potentially leading to enormous stresses on the lithospheres of both bodies. We conclude that if \emph{New Horizons} uncovers evidence for ancient tectonic activity consistent with despinning on Pluto and/or tidal recession on Charon, the most self-consistent explanation is that Pluto had an ocean during the time period of Charon's orbital evolution. Despinning stresses on Pluto would manifest as normal faults at the poles, strike-slip features in the mid-latitudes, and equatorial thrust faults \citep{Melosh77}. Stresses on Charon arising from the collapse of its tidal bulge would create normal faults at the poles, strike-slip faults in the mid-latitudes, with a zone of thrust faults at the sub-Pluto point \citep{Melosh80tide}. Concurrent volume changes in the interiors of Pluto or Charon could add to these predicted stress states and significantly shift the boundaries of the tectonic regions toward or away from the poles or tidal axis. If \emph{no} tectonic features are observed, this could indicate an extremely cold Pluto pre- or post-impact (and thus a very long timescale for orbital evolution), or that tectonic features had been removed or buried by other processes. Tectonic features formed by the post-impact evolution could have been obscured by, e.g., later flooding by cryomagmas. It is also possible that infall of Pluto's tenuous atmosphere may subdue landforms created by ancient tectonics. We find that the orbital evolution time scale for Charon depends sensitively on the interior thermal state of Pluto. This is because the tidal quality factor of Pluto, which controls the rate at which angular momentum is transferred from Pluto's spin to Charon's orbit, depends on $Q$, which in turn depends on the viscosity of its interior. A cold, highly viscous (``stiff'') Pluto will shed angular momentum more slowly than a warm, low-viscosity Pluto. If Pluto is differentiated and has a liquid water ocean, its ice shell is no longer rigidly coupled to its rock core, permitting large tidal deformations and rapid tidal evolution even if the ice shell has a high viscosity. For Charon to evolve to the dual synchronous state on a time scale less than the age of the solar system, a fully differentiated Pluto with an ocean (Figure \ref{fig:ocean}) must be warmed post-impact to a temperature $T \gtrsim 170$ to 190 K, so that the viscosity of the convecting portion of its ice shell, $\eta_i < 10^{19}$ Pa s. Similar conditions must be achieved in an undifferentiated Pluto. However, we consider it likely that Pluto is differentiated either before or during Charon's tidal evolution because the energy liberated during accretion is sufficient to melt its ice, Pluto is unstable to runaway differentiation, and its high mean density implies a high rock fraction in the interior and thus, ample radiogenic heat. If Pluto did not have an ocean during Charon's orbital evolution, its interior must be warmed to $\sim 240$ to 260 K for Pluto and Charon to evolve to the dual synchronous state. Temperatures this high would lead to melting, particularly if Pluto has even just a small amount of ammonia or low-eutectic salt mixed in with its ices. While it is possible for orbital evolution to drive tectonics on Pluto and Charon for any of the assumed interior models for Pluto, the thermal/orbital pathway of the system depends sensitively on the interior states of Pluto and Charon, which may change over time. For example, in the differentiated model without an ocean, tectonics can only occur if Charon starts close to its current location, and Pluto's ice is very close to the melting point (but not molten). In the undifferentiated model, tectonic activity occurs if the ice has a low viscosity, but this same low viscosity ice will dissipate enough heat to begin melting and possibly trigger runaway differentiation. In the differentiated Pluto model with an interior ocean, tectonic activity occurs above an interior that has the right viscosity for a floating, conducting or convecting ice shell. Thus, more detailed calculations (e.g., \citealt{RobuchonNimmo}) will be warranted once the \emph{New Horizons} mission has returned its data. | 14 | 3 | 1403.6377 |
1403 | 1403.7147_arXiv.txt | If binary black holes form following the successive core collapses of sufficiently massive binary stars, precessional dynamics may align their spins $\mathbf S_1$ and $\mathbf S_2$ and the orbital angular momentum $\mathbf L$ into a plane in which they jointly precess about the total angular momentum $\mathbf J$. These spin orientations are known as spin-orbit resonances since $\mathbf S_1$, $\mathbf S_2$, and $\mathbf L$ all precess at the same frequency to maintain their planar configuration. Two families of such spin-orbit resonances exist, differentiated by whether the components of the two spins in the orbital plane are either aligned or antialigned. The fraction of binary black holes in each family is determined by the stellar evolution of their progenitors, so if gravitational-wave detectors could measure this fraction they could provide important insights into astrophysical formation scenarios for binary black holes. In this paper, we show that even under the conservative assumption that binary black holes are observed along the direction of $\mathbf J$ (where precession-induced modulations to the gravitational waveforms are minimized), the waveforms of many members of each resonant family can be distinguished from all members of the other family in events with signal-to-noise ratios $\rho \simeq 10$, typical of those expected for the first detections with Advanced LIGO/Virgo. We hope that our preliminary findings inspire a greater appreciation of the capability of gravitational-wave detectors to constrain stellar astrophysics and stimulate further studies of the distinguishability of spin-orbit resonant families in more expanded regions of binary black-hole parameter space. | Gravitational waves (GWs) emitted during the inspiral of binary black holes (BBHs) are expected to be an important source \cite{2010CQGra..27q3001A} for future networks of GW detectors such the Advanced Laser Interferometer Gravitational Wave Observatory (LIGO) and Virgo \cite{2010CQGra..27h4006H}, LIGO-India \cite{2013IJMPD..2241010U}, the Kamioka Gravitational Wave Detector (KAGRA) \cite{2012CQGra..29l4007S}, and the Einstein Telescope \cite{2010CQGra..27s4002P}. These BBHs can form in two distinct channels: (1) mass segregation can cause isolated black holes (BHs) to sink to the centers of dense stellar clusters and dynamically form binaries \cite{2006ApJ...637..937O,2007PhRvD..76f1504O}, or (2) massive binary stars can evolve into BBHs if each member of the binary is sufficiently massive at the time of core collapse and binary evolution does not destroy the binary before both stars have had the chance to collapse into BHs \cite{2012ApJ...759...52D,2013ApJ...779...72D}. Once formed, BBHs emit GWs that extract energy and angular momentum from the orbit, decreasing the binary separation and increasing the orbital frequency (and thus the GW frequency). Most binaries are expected to circularize by the time they enter the sensitivity band of ground-based detectors \cite{1963PhRv..131..435P,1964PhRv..136.1224P} (see \cite{2013PhRvD..87d3004E} and references therein for recent work on eccentric binary rates and detection strategies). Circular BBH inspirals are characterized by eight intrinsic parameters: the masses $m_1$ and $m_2$ of each BH and their spins $\mathbf S_1$ and $\mathbf S_2$. We choose without loss of generality for the first BH to be more massive than the second: $m_1 > m_2$. The spectrum of emitted GWs depends on these eight parameters, which can therefore be measured by GW detectors if the sources are observed with a sufficient signal-to-noise ratio $\rho$. The distributions of these intrinsic parameters depend on how the BBHs form, allowing GW parameter estimation to constrain not just individual BBH systems, but their astrophysical formation channels as well. \begin{figure} \centering \includegraphics[width=0.9\columnwidth]{fig1} \caption{(color online) Conventions and definitions used in this paper. We work in the radiation frame, where the $z$ axis is oriented along the line of sight $\hat{\mathbf{n}}$. The orbital angular momentum $\mathbf{L}$ lies in the $xz$-plane at $f_{\rm ref}$ and is inclined by an angle $\iota$ with respect to the line of sight. The directions of the spins $\mathbf{S_1}$ (blue) and $\mathbf{S_2}$ (red) are specified using polar angles $\theta_i$ and azimuthal angles $\Phi_i$ $(i=1,\,2)$, which are defined in a frame where the $z$ axis is aligned with the orbital angular momentum $\mathbf{L}$. As resonant binaries precess, their orbital angular momentum and spins remain coplanar implying that the angle $\Delta\Phi=\Phi_2 - \Phi_1$ (green) remains fixed at either $0^\circ$ or $\pm180^\circ$. In later sections of the paper, we will fix $\iota$ and $\Phi_1$ by aligning the line of sight with the total angular momentum: cf. Eqs.~(\ref{Jalongn_phi1}) and (\ref{Jalongn}). } \label{radframe} \end{figure} Our focus in this paper is on whether BBH spin orientations can be measured with sufficient accuracy in $\rho \simeq 10$ sources to constrain the formation of binaries. BBH spin directions are described by three parameters: the two angles $\theta_i$ between spins $\mathbf S_i$ and the orbital angular momentum $\mathbf L$ and the angle $\Delta\Phi = \Phi_2 - \Phi_1$ between the components of the two spins in the orbital plane (see Fig.~\ref{radframe}). Although the individual angles $\Phi_i$ of each of the BBH spins are among the 8 observable intrinsic parameters listed above, only their difference $\Delta\Phi$ provides constraints on BBH formation in the absence of an additional intrinsic vector to break the axisymmetry of the equatorial plane. In the first astrophysical formation channel described above, the BBHs form independently and the dynamical formation of the binary should not depend on the BH spin. We therefore expect both BH spins to have isotropic orientations, in which case the post-Newtonian (PN) GW inspiral will preserve the isotropy of the BBH spins \cite{2007ApJ...661L.147B}. This is not the case, however, in the second astrophysical formation channel, where the BBHs inherit the directions of their spins from their stellar progenitors. In a previous paper \cite{2013PhRvD..87j4028G}, we examined how the spins of BBHs formed from stellar binaries depend on the evolution of their stellar progenitors. Throughout this evolution, the initially more massive star will be designated as the ``primary" and the less massive star will be called the ``secondary". The binary evolution proceeds in several stages: \begin{itemize} \item[(a)] The binary stars initially have spins aligned with their orbital angular momentum $\mathbf L$ as tidal alignment occurs on a much shorter timescale than the main-sequence lifetimes of the stars \cite{2001ApJ...562.1012E}. \item[(b)] The more massive primary evolves more quickly than the secondary, filling its Roche lobe and transferring mass to the secondary. \item[(c)] The core of the primary collapses, forming a BH and detonating a supernova explosion. This asymmetric explosion kicks the binary and tilts the orbital plane. The directions of the stellar spins remain unchanged and thus become misaligned with the new direction of $\mathbf L$. \item[(d)] Tides align the spin of the secondary with the new direction of $\mathbf L$ while leaving the spin of the more compact BH unchanged. \item[(e)] The core of the secondary collapses into a BH. The orbital plane is tilted a second time, misaligning the spin of the secondary with the new direction of $\mathbf L$ and on average increasing the misalignment of the spin of the primary even further. \item[(f)] The BBH spins precess many times before the frequency of emitted GWs enters the sensitivity band of ground-based detectors. \end{itemize} Although the PN spin precession in stage (f) above leaves isotropic spin distributions isotropic as the BBHs inspiral, it can profoundly affect anisotropic spin distributions resulting from stages (a) - (e). The manner in which spin precession alters the distribution of BBH spins can best be understood by appreciating the influence of PN spin-orbit resonances, first identified by Schnittman \cite{2004PhRvD..70l4020S}. BBHs evolve on three distinct timescales: (1) the orbital time $t_{\rm orb} \sim (r^3/GM)^{1/2}$, (2) the precession time $t_{\rm pre} \sim c^2r^{5/2}/[\eta(GM)^{3/2}] \sim (t_{\rm orb}/\eta)(r/r_g)$, and (3) the radiation-reaction time $t_{\rm RR} \sim E/|dE_{\rm GW}/dt| \sim c^5 r^4/[\eta(GM)^3] \sim (t_{\rm orb}/\eta)(r/r_g)^{5/2}$, where $M = m_1 + m_2$ is the total mass, $\eta = m_1m_2/M^2$ is the symmetric mass ratio, and $r_g = GM/c^2$ is the gravitational radius. In the PN regime, $r \gg r_g$ and these timescales are widely separated: $t_{\rm orb} \ll t_{\rm pre} \ll t_{\rm RR}$. In this limit, we can average the spin-precession equations \cite{1995PhRvD..52..821K,2009PhRvD..79j4023A,2011PhRvD..84d9901A} over an orbit while leaving the total angular momentum $\mathbf{J} = \mathbf{L} + \mathbf{S_1} + \mathbf{S_2}$ fixed. The three angular momenta $\mathbf L$, $\mathbf S_1$, and $\mathbf S_2$ will generally span three-dimensional space at any given time and precess in a complicated fashion on the precession time $t_{\rm pre}$ that preserves the magnitude and direction of $\mathbf J$. However, Schnittman discovered special spin configurations in which $\mathbf L$, $\mathbf S_1$, and $\mathbf S_2$ would remain in a two-dimensional plane and jointly precess about $\mathbf J$ on the precession time $t_{\rm pre}$ \cite{2004PhRvD..70l4020S}. He called these configurations ``spin-orbit resonances'' because $\mathbf L$, $\mathbf S_1$, and $\mathbf S_2$ all precessed about $\mathbf J$ at the same frequency. These spin-orbit resonances are divided into two families: resonances in which the spin components in the orbital plane are aligned ($\Delta\Phi = 0^\circ$) and those in which these components are antialigned ($\Delta\Phi = \pm 180^\circ$). At a given binary separation $r$ [or GW frequency $f = \pi^{-1}(GM/r^3)^{1/2}$], each of the two resonant families defines a different curve in the $\theta_1\theta_2$-plane. As resonant BBHs inspiral on the longer radiation-reaction time $t_{\rm RR}$, they remain in spin-orbit resonances though the values of $\theta_i$ vary as the relationship between $\theta_1$ and $\theta_2$ for the spin-orbit resonances is a function of the separation $r$. One might imagine that since the one-parameter spin-orbit resonances constitute a set of measure zero in the three-dimensional parameter space ($\theta_1$, $\theta_2$, $\Delta\Phi$) of spin configurations at any given separation, they are merely a mathematical curiosity of little relevance to astrophysical BBHs. However, BBHs near a spin-orbit resonance will be influenced by its presence, with $\Delta\Phi$ librating about $0^\circ$ or $\pm 180^\circ$ rather than circulating through the full range $\Delta\Phi \in [-180^\circ, +180^\circ]$. Furthermore, as the binary separation decreases an increasing fraction of BBHs will be captured into this librating portion of the parameter space. Which of the two families will be favored by this capture process, the $\Delta\Phi = 0^\circ$ resonances or the $\Delta\Phi = \pm 180^\circ$ resonances? The answer to this question depends on the distribution of $\theta_i$ at large separations. BBHs where the spin of the more massive BH is less misaligned with the orbital angular momentum than that of the less massive BH ($\theta_1 < \theta_2$) will be preferentially attracted to the $\Delta\Phi = 0^\circ$ family of resonances, while BBHs for which $\theta_1 > \theta_2$ will be preferentially attracted to the $\Delta\Phi = \pm 180^\circ$ family \cite{2004PhRvD..70l4020S}. The distribution of $\theta_i$ at large separations is determined by the astrophysics of BBH formation. If the tidal alignment of the secondary's spin in stage (d) above is efficient, the primary's spin will on average be more misaligned with the orbital angular momentum than the secondary's at the start of PN spin precession in stage (f), since its misalignment will have been built up in {\it both} supernova recoils in stages (c) and (e)\footnote{ The second kick is more likely to increase the misalignment between the orbital angular momentum and the spin of the primary because of the greater amount of phase space at larger values of $\theta$ (the Jacobian determinant $\sin\theta$ increases with $\theta$) provided the first tilt is $\lesssim\pi/2$. See \cite{2013PhRvD..87j4028G} for a discussion.} However, the primary star (which is initially more massive and is thus first to collapse into a BH) will not always become the more massive BH. If enough mass is transferred from the primary to the secondary prior to the first core collapse in stage (b), the primary will evolve into the less massive BH. We will refer to this possibility as the reverse-mass-ratio (RMR) scenario \cite{2013PhRvD..87j4028G}. In this case, the more massive BH (evolved from the secondary) will have a less misaligned spin ($\theta_1 < \theta_2$), and the $\Delta\Phi = 0^\circ$ family of resonances will be preferentially populated. Conversely, in the standard-mass-ratio (SMR) scenario where the primary evolves into the more massive BH, it will have a more misaligned spin ($\theta_1 > \theta_2$) and the $\Delta\Phi = \pm 180^\circ$ family of resonances will be favored. In the ``No Tides" scenario where the tidal alignment in stage (d) is ineffective, neither of the resonant families will be favored over the other. Our previous paper \cite{2013PhRvD..87j4028G} showed that for a simplified but not unreasonable toy model of the BBH formation described above, a large fraction of BBHs were librating about the $\Delta\Phi = 0^\circ~(\pm 180^\circ)$ resonances in the RMR (SMR) scenario by the time the GW frequency $f$ approached the frequency at which most of the signal-to-noise ratio (SNR) is accumulated ($\sim 60$ Hz). In this paper, we investigate what SNR $\rho$ is required to distinguish the GWs emitted by BBHs in the two resonant families. A thorough exploration of the full intrinsic and extrinsic parameter space that characterizes BBH waveforms is computationally prohibitive, so we make several mostly conservative assumptions to restrict this parameter space. To facilitate comparison with our previous paper, we fix $m_1 = 7.5~M_\odot$, $m_2 = 6~M_\odot$, and $\chi_i \equiv S_i/m_i^2 = 1$. These values are close to the expected peak of the distribution of astrophysical BH binaries detectable by Advanced LIGO, as predicted by population-synthesis codes \cite{2012ApJ...759...52D}. They are also consistent with the strong influence of PN precession, as binaries are most effectively captured into spin-orbit resonances when the BBH masses are comparable ($q \equiv m_2/m_1 \lesssim 1$) and both dimensionless spin amplitudes are large ($\chi_i \gtrsim 0.5$) \cite{2010PhRvD..81h4054K,2013PhRvD..87j4028G}. All BBHs, therefore, have the same masses and spin magnitudes, ensuring that the spin directions are solely responsible for the differences in the waveforms. We also choose the position $\hat{\mathbf n}$ of the BBHs on the sky such that they are directly overhead of the GW detectors. We align the direction $\hat{\mathbf J}$ of the total angular momentum with $\hat{\mathbf n}$ at a reference frequency $f_{\rm ref} = 60$~Hz, a typical frequency at which most of the SNR is accumulated. This latter choice is conservative since $\hat{\mathbf L}$ precesses about the nearly constant $\hat{\mathbf J}$ during the inspiral, and thus the precessional modulations to the waveform due to changes in the angle between $\hat{\mathbf L}$ and $\hat{\mathbf n}$ are minimized. With these choices, we compare the waveforms of each member of the two families of spin-orbit resonances with those of all of the members of the opposite family by computing their overlap ${\cal O}$. If this overlap with all members of the opposite family is sufficiently less than unity, we can safely claim to have determined to which of the resonant families the BBH belongs. The remainder of this paper is organized as follows. In Sec.~\ref{sec:Dynamics} we review the dynamics of resonant BBHs, introduce a convenient parametrization to identify members of each resonant family, and show qualitatively why the two families are dynamically distinguishable. In Sec.~\ref{sec:Compare}, we examine the GWs emitted by resonant BBHs, use the overlap between waveforms from different families to assess their distinguishability, then investigate how this distinguishability can be used to differentiate between astrophysical scenarios of BBH formation. In Sec.~\ref{sec:SingleSpin} we hypothesize that the dynamics and waveforms of resonant binaries are similar to binaries with a single effective spin, then use this hypothesis to develop two different predictions for the best matching waveforms in the different resonant families. Some final remarks are provided in Sec.~\ref{S:disc}. Some technical details concerning the numerical evolution of the BBHs and the nature of the correspondence between matching waveforms in the two families are given in Appendixes \ref{app_coord} and \ref{ap:Symmetry}. Throughout the rest of this paper we use geometrical units where $G=c=1,$ and we use hats to identify unit vectors. For example, the direction of the orbital angular momentum will be denoted by $\hat{\mathbf L}={\mathbf L}/|{\mathbf L}|$. | \label{S:disc} BBH formation remains shrouded in mystery. Such systems are predicted to be very rare; none have been observed to date, which is not surprising, given their minimal electromagnetic signature. BBH mergers are copious sources of GWs, however, so they should be a prominent signal for GW detectors, in contrast to electromagnetic telescopes. GW detectors can, in principle, measure all of the intrinsic parameters associated with a binary if that binary is detected with a sufficient SNR $\rho$. Our previous paper \cite{2013PhRvD..87j4028G} established a surprisingly tight connection between BBH spin orientations and BBH formation: binaries with an efficient tidal alignment that undergo a mass-ratio reversal will preferentially be found in the $\Delta\Phi = 0^\circ$ family of resonances, those that fail to undergo such a reversal will preferentially be found in the $\Delta\Phi = \pm180^\circ$ family of resonances, and those without an efficient tidal alignment are equally likely to be found in either resonant family. A measurement of the fraction of BBHs in each resonant family could therefore be used to distinguish between different astrophysical scenarios of BBH formation. This paper is the first attempt to assess the feasibility of such a proposed measurement. The qualitatively distinct spin orientations in the two families lead to quantitative differences in the amount of orbital-plane precession. The greater misalignment between the orbital angular momentum $\mathbf L$ and the total angular momentum $\mathbf J$ in the $\Delta\Phi = 0^\circ$ family implies greater precessional modulation of the resulting waveforms, even under the conservative assumption that binaries are viewed from a direction $\hat{\mathbf{n}} = \hat{\mathbf{J}}$ where precessional modulation is minimized. Precession-induced differences between the waveforms generated by binaries in the two resonant families lead to a maximum overlap ${\cal O}_{\rm max}(\xi_{\rm source}) < 1$ between a source with projected effective spin $\xi_{\rm source}$ in one family and the best matching template $\xi_{\rm template}^{\rm BM}$ from the other family. The slow variation of ${\cal O}_{\rm max}(\xi_{\rm source})$ implies that this matching is symmetric to better than a part in $10^3$: the binary from the first family with $\xi_{\rm source}$ is also very nearly the template that provides the best match when the binary from the second family with $\xi_{\rm template}^{\rm BM}$ is serving as the source. The resonant family of a binary with $\xi_{\rm source}$ can be identified when ${\cal O}_{\rm max}(\xi_{\rm source}) < 1- \rho^{-2}$; this condition holds for much of our one-parameter space $\xi_{\rm source} \in [-1, +1]$ for $\rho \gtrsim 10$, a typical SNR expected for the first GW detections.\footnote{ After our study was completed, the authors became aware of a work by Vitale et {\it al.}~\cite{2014arXiv1403.0129V} that performs detailed parameter estimation on selected generic double-spin binaries. Unfortunately, these authors did not select resonant configurations for their detailed investigation (even if their injected configurations are coplanar at $f=100~{\rm Hz}$). } Different astrophysical BBH formation scenarios can be distinguished if they predict that measurably different fractions of binaries reside in the portions of parameter space that can be identified by the criterion above as belonging to each of the resonant families. This is indeed the case for the three scenarios described in our previous paper \cite{2013PhRvD..87j4028G}; if 100 binaries are detected with $\rho \gtrsim 10$, $\sim15$ should be found in the $\Delta\Phi = 0^\circ$ family in the reverse-mass-ratio scenario, $\sim20$ should be found in the $\Delta\Phi = \pm180^\circ$ family in the standard-mass-ratio scenario, and $\sim5$ should be found in {\it each} family if the tidal alignment is inefficient. These three scenarios and the resulting distributions of BBH spin orientations were constructed long before we calculated our first overlap, and thus are in no way optimized to maximize the number of binaries in the identifiable portion of parameter space. Finally, except for contrived scenarios, BBHs should be detected frequently \cite{2010CQGra..27q3001A,2012ApJ...759...52D}, with a rate of events at SNR $>\rho$ roughly proportional to $\simeq O(1-1000)\unit{yr}^{-1}(10/\rho)^3$ at the design sensitivity. Extrapolating from our results, only for pessimistic scenarios do we expect to have too few and too faint events to distinguish between the RMR and SMR scenarios. Our claim that GW detectors can be used to constrain BBH formation scenarios must remain provisional until more realistic higher-dimensional model parameter spaces are considered.\footnote{A recent paper \cite{2014CQGra..31j5017G} has some overlap with our own and also argues that the two gravitational-wave signals can be distinguished.} Our demonstration that the single-spin approximation describes resonant binaries with reasonable accuracy may facilitate such a higher-dimensional analysis, but this remains a subject for future work. Our current study offers the tantalizing promise that Advanced LIGO/Virgo may not only discover GWs and test general relativity in the strong-field regime, but also may revolutionize our understanding of astrophysical BBH formation. | 14 | 3 | 1403.7147 |
1403 | 1403.7976_arXiv.txt | {The determination of the magnetic field vector in quiescent solar prominences is possible by interpreting the Hanle and Zeeman effects in spectral lines. However, observational measurements are scarce and lack high spatial resolution.} {To determine the magnetic field vector configuration along a quiescent solar prominence by interpreting spectropolarimetric measurements in the \ion{He}{i}~1083.0~nm triplet obtained with the Tenerife Infrared Polarimeter installed at the German Vacuum Tower Telescope of the Observatorio del Teide.} {The \ion{He}{i}~1083.0~nm triplet Stokes profiles are analyzed with an inversion code that takes into account the physics responsible of the polarization signals in this triplet. The results are put into a solar context with the help of extreme ultraviolet observations taken with the Solar Dynamic Observatory and the Solar Terrestrial Relations Observatory satellites.} {For the most probable magnetic field vector configuration, the analysis depicts a mean field strength of 7~gauss. We do not find local variations in the field strength except that the field is, in average, lower in the prominence body than in the prominence feet, where the field strength reaches $\sim$~25~gauss. The averaged magnetic field inclination with respect to the local vertical is $\sim$~77\degr. The acute angle of the magnetic field vector with the prominence main axis is 24\degr\/ for the sinistral chirality case and 58\degr\/ for the dextral chirality. These inferences are in rough agreement with previous results obtained from the analysis of data acquired with lower spatial resolutions.} {} | \label{sec:intro} Although the first photographic plates of prominences were taken more than 150 years ago it took 100 years to discover, by means of the first spectropolarimetric measurements in prominences, that these solar structures are clear manifestations of the confinement of plasma within giant magnetic structures\footnote{There are many papers about solar prominences, but we recommend the reader to go first through the historical work of Einar Tandberg-Hanssen (e.g., \citealt{1998ASPC..150...11T,2011SoPh..269..237T}).}. Prominences, also referred to as filaments when observed against the solar disk, are cool, dense, magnetized formations of 10$^4$~K plasma embedded into the 10$^6$~K solar corona (for reviews see \citealt{2010SSRv..151..333M,2010SSRv..151..243L}). They are located above Polarity Inversion Lines (PILs or filament channels), i.e., the line that divides regions of opposite magnetic flux in the photosphere. Morphologically speaking, prominences can be separated in different classes \citep{1943ApJ....98....6P,1995ASSL..199.....T}. Among them, quiescent prominences are seen as large sheet like structures suspended above the solar surface and against gravity. The global structure of quiescent prominences change little with time, preserving their global shape during days and even weeks. Locally, they consist of fine and vertically oriented plasma structures, so-called threads, that evolve continually (e.g., \citealt{1976SoPh...49..283E,1994SoPh..150...81Z}). Recent observations taken with the Hinode satellite have revolutionized our knowledge of quiescent prominences fine scale structuring and dynamics; for instance, plasma oscillations, supersonic down-flows, or plasma instabilities like the Rayleigh-Taylor instability in prominence bubbles \citep{2008ApJ...676L..89B,2010ApJ...716.1288B,2008ApJ...689L..73C,2007Sci...318.1577O}. The magnetic configuration of quiescent prominences has been investigated by using first the longitudinal Zeeman effect and later by measuring the full Stokes vector in spectral lines sensitive to the joint action of the Zeeman and Hanle effect (e.g., \citealt{1989ASSL..150...77L} and \citealt{2007ASPC..368..291L} for reviews). For instance, spectropolarimetric observations in the \ion{He}{I}~D3 multiplet at 587.6~nm have greatly contributed to the understanding of the magnetic field configuration in prominences \citep{1983SoPh...89....3A,1985SoPh...96..277Q,2003ApJ...598L..67C}. Full Stokes polarimetry in the \ion{He}{I}~D3 multiplet at 587.6~nm is accessible from several ground-based observatories, such as the French-Italian telescope (THEMIS) at the Observatorio del Teide, the Advanced Stokes Polarimeter at the Dunn Solar Telescope in Sacramento Peak, or the Istituto Ricerche Solari (IRSOL) observatory. Of particular interest to infer the magnetic field vector in prominences is the \ion{He}{I} triplet at 1083.0~nm. This spectral line can be clearly seen in emission in off-limb prominences (e.g., \citealt{2006ApJ...642..554M}) and in absorption in on-disk filaments (e.g., \citealt{1998ApJ...493..978L}). The \ion{He}{I} triplet is sensitive to the joint action of atomic level polarization (i.e., population imbalances and quantum coherence among the level's sublevels, generated by anisotropic radiation pumping) and the Hanle (modification of the atomic level polarization due to the presence of a magnetic field) and Zeeman effects \citep{2002Natur.415..403T,2007ApJ...655..642T}. This fact makes the \ion{He}{I}~1083.0~nm triplet sensitive to a wide range of field strengths from dG (Hanle) to kG (Zeeman). Importantly, the \ion{He}{I} 1083.0~nm triplet is easily observable with the Tenerife Infrared Polarimeter (TIP-II; \citealt{2007ASPC..368..611C}) installed at the German Vacuum Tower Telescope (VTT) of the Observatorio del Teide (Tenerife, Spain). Finally, an user-friendly diagnostic tool called ``HAZEL'' (from HAnle and ZEeman Light) is available for modeling and interpreting the \ion{He}{I}~1083.0~nm triplet polarization signals, easing the determination of the strength, inclination and azimuth of the magnetic field vector in many solar structures \citep{2008ApJ...683..542A}. The HAZEL code has already been used to analyze \ion{He}{I} 1083.0~nm triplet spectropolarimetric data of prominences \citep{2013hsa7.conf..786O}, spicules \citep{2010ApJ...708.1579C,2012ApJ...759...16M}, sunspot's super-penumbral fibrils \citep{2013ApJ...768..111S}, emerging flux regions \citep{2010MmSAI..81..625A}, and the quiet solar chromosphere \citep{2009ASPC..405..281A}. The HAZEL code has also been applied to \ion{He}{I}~D3 observations of prominences and spicules \citep{2011ASPC..437..109R}. However, the information we have about the spatial variations of the magnetic field vector in solar prominences is still very limited because of the insufficient spatial resolution of the observations, restricted to single point measurements (e.g., \citealt{1983SoPh...83..135L,1983SoPh...89....3A}), single slit measurements (e.g., \citealt{2006ApJ...642..554M}), or two-dimensional slit scans (e.g., \citealt{2003ApJ...598L..67C,2007ASPC..368..347M}) at spatial resolutions of about 2\arcsec, much lower than the sub-arcseconds resolutions achieved by the Hinode spacecraft in prominence broad-band imaging. We have an approximate picture of the global magnetic properties of quiescent solar prominences, mainly thanks to the information encoded in spectral lines sensitive to the Hanle and Zeeman effect. The magnetic field in quiescent prominences is rather uniform and has mean field strengths of tens of gauss, typically in the range of 3~G to 30~G. The magnetic field vector forms an acute angle of about 35\degr\/ with the prominence long axis \citep{1970SoPh...15..158T,1983SoPh...83..135L,1994SoPh..154..231B,2003ApJ...598L..67C}. The field lines are found to be highly inclined with respect to the local vertical (e.g., \citealt{1983SoPh...89....3A}). For instance, \cite{1983SoPh...83..135L} found a mean inclination of 60\degr\/ from the local vertical in a sampling of 15 prominences. Their data were limited to single point measurements and their estimated rms error was about 15\degr. More recently, \cite{2003ApJ...598L..67C,2005ApJ...622.1265C} inferred the vector field map in a quiescent prominence and found inclinations of about 90\degr\/ with respect to the local vertical. These authors also reported that the field can be organized in patches where it increases locally up to 80~G. The magnetic configuration seems to be different for polar crown prominences where the field is found to be inclined by about 25\degr\/ with respect to the solar radius vector through the observed point \citep{2006ApJ...642..554M}. Finally, it has been found that, for 75\% of the analyzed prominences, the perpendicular component of the magnetic field vector to the prominence long axis or PIL points to the opposite direction with respect to the photospheric magnetic field. In this case, they are classified as inverse polarity prominences \citep{1983SoPh...83..135L}. \begin{figure}[!t] \begin{center} \resizebox{\hsize}{!}{\includegraphics{fig1.eps}} \end{center} \caption{Peak intensity map of the \ion{He}{I}~1083.0~nm triplet emission profile. The prominence is seen as a bright structure against a dark background. The bottom, dark part corresponds to the solar limb. The top-right arrow points to the solar North direction. The de-projected height (see Sect.\ref{sec2}) above the solar surface is shown on the right axis. The data was taken on 20 May 2011, at 9:44 UT and finished at 11:15 UT, within the same day.} \label{fig1} \end{figure} \begin{figure*}[!t] \begin{center} \resizebox{0.9\hsize}{!}{\includegraphics{fig2.eps}} \end{center} \caption{Illustrations of the observed prominence as seen in SDO/AIA, STEREO/EUVI, and the BBSO. Top panels correspond to \ion{Fe}{IX}~171~\AA, \ion{Fe}{XIV}~211~\AA\/ and \ion{He}{II}~304~\AA\/ AIA band-pass filter images. Bottom panels are BBSO H$\alpha$ broad-band image, and \ion{Fe}{XII} 195 \AA\/ and \ion{He}{II} 304 \AA\/ STEREO-B/EUVI band-pass images. All images were taken on 20 May 2011, the same day the TIP-II observations were carried out. The white box represents the TIP-II field-of-view. Lines of constant Stonyhurst heliographic longitude and latitude on the solar disk are overplotted. Axis are in heliocentric coordinates. The arrows pinpoint the location of horn-like structures in \ion{He}{II} 304 \AA. The TIP-II slit virtual position can be seen in the botton central panel. The quiescent prominence can be clearly seen in the AIA images as well as in H$_\alpha$. In STEREO-B, it can be seen as a dark, elongated structure. The temporal evolution of the prominence in the SDO/AIA \ion{Fe}{XIV}~211~\AA\/ channel is available in the on-line edition.} \label{fig2} \end{figure*} The magnetic field vector we infer through the interpretation of polarizations signals, such as those of the \ion{He}{I} 1083.0~nm multiplet, is associated with the coolest and densest prominence material. For this reason, the magnetic field in prominences has also been investigated by indirect means, i.e., constructing models of the field geometry in order to capture the observed prominence shape and properties \citep{1998A&A...329.1125A,1999A&A...342..867A,1998A&A...335..309A,2012ApJ...761....9D}. Such studies have contributed to our present picture of the global magnetic field structure associated to the prominence, although most models assume that the prominence material is suspended in magnetic dips. Among these models, we have the sheared-arcade models \citep{1989ApJ...343..971V} and the twisted flux-rope models \citep{1994SoPh..155...69R}. In the first ones, an helical magnetic structure is generated via photospheric shear flow motions that give rise to magnetic reconnection of pre-existing magnetic fields lines near the PIL. The cool material of the prominence is then supported by the magnetic dips of the helical structure via a magnetic tension force \citep{1957ZA.....43...36K,2005ApJ...626..551L,2010ApJ...714..618C}, by MHD-waves pressure \citep{2000SoPh..194...73P}, or by the presence of tangled magnetic fields in very small scales \citep{2010ApJ...711..164V}. The twisted flux-rope models suggest that the helical magnetic field structure supporting the prominence material has emerged from below the photosphere. Both models yield magnetic properties compatible with the present, low-resolution observational constraints. The local magnetic field in prominences has also been investigated by interpreting the dynamics of rising plumes using magnetohydrodynamic models \citep{2012ApJ...746..120H,2012ApJ...761..106H}. Here we present the results of the analysis of ground-based spectropolarimetric observations of the \ion{He}{I}~1083.0~nm triplet taken in a quiescent solar prominence. The data were obtained with the TIP-II instrument \citep{2007ASPC..368..611C} installed at the German VTT at the Observatorio del Teide. This instrument is providing observations of solar prominences at spatial resolutions of about 1\arcsec--1\farcs5 during regular observing conditions, and even below one arcsecond in periods of excellent seeing conditions. In this paper, some of these new observations will be analyzed with the HAZEL code. We will first describe the observations (sections 2 and 3) and then explain the diagnostic technique (section 4). In section 5 we present the inferred two dimensional map of the magnetic field vector of the observed prominence, and then we discuss and summarize the results in section 6. | \label{sec5} In this paper we have shown that spectropolarimetric observations in the \ion{He}{i}~1083.0~nm triplet are very useful to determine the strength and orientation of the magnetic field in solar prominences. In particular, we have determined the magnetic field vector in a quiescent solar prominence, using observations of the \ion{He}{i}~1083.0~nm triplet taken with the TIP-II instrument installed at the VTT. Even though the integration time per slit position was of the order of one minute and the total time to scan the prominence took $\sim$~1.5 hour, it is possible to acquire data under stable observing conditions with TIP-II and still maintaining a moderate spatial resolution of 1\arcsec--1\farcs5. These integration times are necessary to increase the signal-to-the-noise ratio and detect both, circular and linear polarization signals in the prominence. These signals are often buried in the noise in ``standard'' TIP-II observations. The detection of the linear polarization signals allows us to determine the orientation of the field vector while Stokes V is crucial to fix the field strength. To infer the field vector we have employed the HAZEL inversion code, which includes all necessary physics for interpreting the Stokes I, Q, U, and V profiles. We have shown that the use of context data, such as that provided by STEREO and SDO, may be crucial for setting up the scattering problem. In our case, STEREO allowed us to determine the prominence heights, its position on the solar disk (viewing angle with respect to the solar vertical), and the orientation of the PIL with respect to the solar limb (or to the LOS). The \ion{He}{i}~1083.0~nm triplet suffers from two ambiguities: the 90\degr\/ ambiguity and the 180\degr\/ ambiguity of the Hanle effect. We have shown that the observed profiles themselves do not encode sufficient information to solve any of these ambiguities in 90\degr\/ scattering geometry (see Fig.~3). Fortunately, from theoretical arguments we can discard two of the solutions. For instance, we can assume that the magnetic field vector belongs to a weakly twisted flux rope or to a sheared arcade which also contains weakly twisted field lines. In this case, the prominence material would be located in dipped magnetic field lines. In these two models the component of the magnetic field vector along the prominence main axis dominates. Thus, the quasi-vertical solutions, which suggest that the magnetic field vector is almost perpendicular to the prominence main axis can be discarded. The quasi-vertical solution would also require a extremely highly twisted structure, which is rather improbable in prominences \citep{1989ApJ...343..971V}. Thus, there are two possible solutions only where the magnetic field vector is highly inclined, the quasi-horizontal solutions. These two solutions are connected through the 180\degr\/ ambiguity of the Hanle effect, i.e., they differ by a 180\degr\/ rotation in the plane perpendicular to the LOS. The quasi-horizontal solutions confirm previous findings about the average field strength in quiescent prominences, being about 7~G, in average. An interesting result is that the field strength seems to be more intense at the prominence feet, reaching values up to 30~G and coinciding with areas where the opacity increases. If the dense plasma is truly suspended in magnetic dips, the existing correlation between the opacity and the field strength may provide additional information to understand current physical mechanisms for suspending the prominence material. Interestingly we do not detect abrupt changes in the prominence field strength contrary to the results of \cite{2003ApJ...598L..67C}. Our results for the orientation of the field vector with respect to the solar surface slightly deviate from previous measurements. In particular, we found that the field vector is about 77\degr\/ inclined with respect to the solar vertical. This result is between the values reported by e.g., \cite{1989ASSL..150...77L} and \cite{1994SoPh..154..231B}, with inclinations of about 60\degr\/ from the local vertical, and those reported by \cite{2003ApJ...598L..67C}, mostly horizontal fields. Regarding the orientation of the field with respect the prominence main axis, we found it to be $\sim$~58\degr\/ in one case (dextral chirality) and $\sim$~156\degr\/ in the corresponding 180\degr\/ ambiguous solution (sinistral chirality). The first one differs from previous findings. For instance, \cite{1989ASSL..150...77L}, \cite{1994SoPh..154..231B}, and \cite{2003ApJ...598L..67C} report angles below 30\degr\/. In contrast, the second solution, $\chi^\dag=156^\circ$, implies an acute angle of $24^\circ$, in line with previous measurements. In practice, we cannot distinguish between the two of them. We point out that the sinistral chirality case is the most probable solution for southern prominences \citep{martin}, although it may also be possible that the twisting is related with the fact that this prominence was ejected few hours after the observations. In this case, the more twisted case (dextral) would be the ``true'' magnetic configuration. | 14 | 3 | 1403.7976 |
1403 | 1403.0621_arXiv.txt | The main aim of this paper is to study the astrophysical behaviour of open clusters' properties along the Milky Way Galaxy. Near-IR {\it JHK$_{S}$} {\it (2MASS)} photometry has been used for getting a homogeneous Catalog of 263 open clusters' parameters, which are randomly selected and studied by the author through the last five years; most of them were studied for the first time. The correlations between the astrophysical parameters of these clusters have been achieved by morphological way and compared with the most recent works. | Open star clusters are very important objects in solving problems of star formation, stellar evolution, and improving our knowledge about the distance scale and the kinematic properties of the Milky Way Galaxy. This kind of study requires a large set of homogeneous data on the positions and ages of open star clusters, which are estimated in a precision way from their Colour-Magnitude Diagrams {\it (CMDs)}. However, it is useful to re-investigate the properties and structures of the Milky Way Galaxy using the most recent Near--IR {\it JHK$_{S}$} photometric data from the 2--Micron All Sky Survey $(\it {2MASS})$ Point Source Catalogue of Skrutskie et al. (2006). In this context, we used a sample of 263 open star clusters (most of them are studied for the first time) that have been analysed by the author, in a series of papers; Tadross (2008 $\sim$ 2012); see Tables 1 and 2. Our aim is to repeat the work we have done 12 years ago, Tadross 2001 and Tadross et al. 2002, to study such relations using NIR observations instead of UBV. This paper is organized as follows. In Sec. 2, a historical review of the present study is obtained. In Sec. 3, data analysis of the clusters under investigation are presented. The limiting and core radii are given in Sec. 4. The main photometric parameters are obtained in Sec. 5. Ages and locations, distribution are presented in Sec. 6. Reddening distribution is presented in Sec. 7. The diameters and ages' relations are given in Secs. 8 and 9 respectively. Secs. 10 and 11 describe the spiral arms and warp of the Galaxy respectively. The conclusions are obtained in Sec. 12. | The results of our studied clusters in the last five years using near--infrared {\it JHK$_{S}$} photometric system are obtained here, and the correlations between the astrophysical parameters along the Milky Way Galaxy are achieved. It is obvious that ({\it JHK$_{S}$}) 2MASS system affected the magnitude limit of the clusters, which detects many faint members located away from the cluster's core, so then the cluster seems to be larger than in optical bands. Detecting stars located in the lower parts of CMDs, make the fitting with standard zero age main sequence much easier. This, of course, has contributed to the evaluation of the cluster parameters, i.e. distances, diameters, ages, reddening, etc. From our reduction, we concluded that $R_{lim}$= 6.85 $R_{c}$; for the clusters up to $R_{c}$= 0.5 arcmin, and $R_{lim}$= 2.88 $R_{c}$; for the clusters up to $R_{c}$= 1.0 arcmin, which are in agreement with Maciejewski \& Niedzielski (2007). We also noticed that the linear size of open clusters increases with ages. The reddening decreases outward the Galactic plane, $Z$ and the Galactic Center, $R_ {gc} $, as well. This is noticed also for clusters located near the Sun vicinity and further than 8.5 kpc from the Galactic centre, i.e. the density of dust and gas decreases, too. \\ \\ From our analysis, we noticed that the number of clusters decreases with $Z$; more than half of the studied clusters ($52\%$) have aged less than 500 mega years and located at average $|Z|$ = 75 pc. Hence, the older ones are located at average $|Z|$ = 275 pc, which is in agreement with Bukowiecki et al. (2011). We can show that the difference between younger and older clusters can be declared in locations and sizes as the following relation: \begin{center} Diam. = 0.53 $R_{gc}$ - 0.19 = 3.18 \emph{Log (age)} - 18.53 \end{center} We found that the number of older clusters increases with $R_{gc}$ and younger ones are obtained at an average $R_{gc} = 8.8$ kpc, which is confirmed by Tadross et al. (2002), Froebrich (2010), and Bukowiecki et al. (2011). The paucity of the clusters at G. longitudes range from $140^{o}$ to $200^{o}$ is noticeable by Tadross et al. (2002), Benjamin (2008), Froebrich (2010), and Bukowiecki et al. (2011). It may reflect the real spatial structure of the Milky Way Galaxy in that direction near the feature region of the Perseus arm (the external youngest arm of the Galaxy). | 14 | 3 | 1403.0621 |
1403 | 1403.2909_arXiv.txt | We have designed, fabricated and characterized absorptive thermal blocking filters for cryogenic microwave applications. The transmission line filter's input characteristic impedance is designed to match $50\,\Omega$ and its response has been validated from 0-to-50\,GHz. The observed return loss in the 0-to-20\,GHz design band is greater than $20\,$dB and shows graceful degradation with frequency. Design considerations and equations are provided that enable this approach to be scaled and modified for use in other applications. | Introduction} Thermal blocking filters find wide use in cryogenic applications ranging from quantum computing to ultra-low-noise detectors. They can be used to provide the environmental isolation between cooled devices and the warmer temperature supporting bias and readout circuitry. In particular, they are effective in rejecting thermal radiation, limiting radio frequency interference, providing a convenient means of heat sinking signal lines, and realizing a vacuum feedthrough. In a microwave instrumentation setting a well-defined characteristic impedance, typically approximating a short or matched boundary condition, is desirable. From a radiometric perspective, such filter structures limit the available power by modifying the transmitted response and effectively reducing the photon density of states in the Planck distribution to a single dimension. A variety of thermal blocking filter construction techniques and designs have been discussed in the literature. In the device's most basic form, a large shunt capacitor forms a single-pole low-pass-filter.~\cite{Bladh2003} More generally, multiple low-pass lumped element stages can be combined in series to produce compact and broadband non-dissipative filter structures.~\cite{Vion1995,Sueur2006,UYen2008} The challenges presented by these implementations include controlling inter-stage isolation and spurious transmission resonances, limiting the filter's total shunt capacitance, and achieving adequate control over circuit parameters as a function of temperature.~\cite{Brown2012} Dissipative solutions based on distributed lossy microwave structures~\cite{Schiffres1964,Martinis1987,Zorin1995,Fukushima1997,Leong2002,Lukashenko2008} can be used to achieve a broadband low-pass transmission response. More recent efforts have strived to retain these desirable properties while providing a well defined impedance match. Specific examples in this class include coaxial-lines~\cite{Milliken2007} and strip-line~\cite{Santavicca2008,Slichter2009} powder filters. In this work, simple matched filter designs based on easily realized absorptive dielectric transmission lines are improved upon, and the resulting performance is described in detail. The filter's response is calculable, repeatable under cryogenic cycling, and is capable of providing an intrinsically broadband matched impedance termination. In Section~\ref{sec:design}, practical design considerations and the governing equations for the filter's operation are described in detail. The fabrication and test of the representative filters are summarized in Section~\ref{sec:fabrication}. | Conclusion} The electrical design, fabrication, and characterization of thermal blocking filters based upon a lossy transmission line are described. These impedance matched structures have been realized in circular and square tube geometries with low manufacturing complexity and terminated with readily available commercial connectors. A return loss $>20$dB from 0-to-20~GHz is observed. A prototype wide slot filter configuration for use with a multi-line fanout board is also presented. The circuit concepts described are flexible and can be readily adapted to meet a host of microwave metrology needs at cryogenic temperatures. \vspace{0pt} | 14 | 3 | 1403.2909 |
1403 | 1403.4856_arXiv.txt | We study the origin of $^{3}$He-rich solar energetic particles ($<$1~MeV\,nucleon$^{-1}$) that are observed consecutively on {\sl STEREO-B}, {\sl ACE}, and {\sl STEREO-A} spacecraft when they are separated in heliolongitude by more than 90\degr. The $^{3}$He-rich period on {\sl STEREO-B} and {\sl STEREO-A} commences on 2011 July 1 and 2011 July 16, respectively. The {\sl ACE} $^{3}$He-rich period consists of two sub-events starting on 2011 July 7 and 2011 July 9. We associate the {\sl STEREO-B} July 1 and {\sl ACE} July 7 $^{3}$He-rich events with the same sizeable active region producing X-ray flares accompanied by prompt electron events, when it was near the west solar limb as seen from the respective spacecraft. The {\sl ACE} July 9 and {\sl STEREO-A} July 16 events were dispersionless with enormous $^{3}$He enrichment, lacking solar energetic electrons and occurring in corotating interaction regions. We associate these events with a small, recently emerged active region near the border of a low-latitude coronal hole that produced numerous jet-like emissions temporally correlated with type III radio bursts. For the first time we present observations of 1) solar regions with long-lasting conditions for $^{3}$He acceleration and 2) solar energetic $^{3}$He that is temporary confined/re-accelerated in interplanetary space. | Solar energetic particle (SEP) $^{3}$He-rich events are characterized by huge abundance enhancements of the rare isotope $^{3}$He (up to factors of $>$10$^{4}$) over solar system abundances. The enrichment of the $^{3}$He is believed to be caused by selective heating due to wave-particle interaction in the flare plasma because of its unique charge to mass ratio \citep[see e.g., reviews by][]{koc84,rea90}. Earlier studies have shown that $^{3}$He-rich SEP events are associated with 2-100~keV electron events \citep{rea85} and their related type III radio bursts \citep{rea86}, which are excited by electrons streaming outward through interplanetary space \citep[e.g.,][]{lin74}. The events have been associated with minor soft X-ray flares \citep[e.g.,][]{zwi78,kah87,rea88}. Recently the sources of $^{3}$He-rich events have been investigated using extreme ultraviolet (EUV) and X-ray solar images \citep{nit06,nit08,wan06}. The responsible source has often been attributed to a jet-like ejection near a coronal hole. Sequences of $^{3}$He-rich SEP events from the same active region (AR) have been observed with a single spacecraft (s/c) \citep{rea86,mas99,mas00,wan06,pic06} only in a limited time interval (about one day) presumably due to loss of the magnetic connection to the flare site. These recurrent events may suggest almost continuous acceleration of $^{3}$He in the single AR \citep{pic06}. In addition, multi-day periods of solar energetic $^{3}$He have been recently discovered with the {\sl Advanced Composition Explorer} ({\sl ACE}) often with no individual $^{3}$He-rich events resolved \citep[and references therein]{mas07}. Such long periods have been interpreted as a combination of continuous production of $^{3}$He-rich SEPs and their confinement in interplanetary magnetic field structures \citep{koc08}. In this paper we report, for the first time, on a $^{3}$He-rich period of SEPs consecutively observed by three widely, in longitude, separated spacecraft {\sl STEREO-B}, {\sl ACE} and {\sl STEREO-A} with delays corresponding to the Carrington rotation rate. The ultimate goal of this study is to identify activity on the Sun and the conditions in interplanetary space leading to the long persistence of energetic $^{3}$He in the heliosphere. In Section~\ref{epo}, we describe energetic particle and solar wind plasma observations on each s/c. The responsible solar sources are identified in Section~\ref{ss} using a combination of coronal field model extrapolations, EUV and radio observations. The results are summarized and discussed in Section~\ref{sad}. | \label{sad} We have examined $^{3}$He-rich periods of SEPs observed consecutively by the {\sl STEREO-B}, {\sl ACE}, and {\sl STEREO-A} spacecraft when they were widely separated in longitude. The period observed by {\sl ACE} consists of two distinct $^{3}$He-rich events. The period observed by {\sl STEREO-B} on 2011 July 1 and later by {\sl ACE} on 2011 July 7 was associated with a sizeable active region, AR 11244. The AR produced energetic electron events with associated type III bursts, soft X-ray flares and a H$\alpha$-flare. The period is characterized by a moderate $^{3}$He-enrichment ($^{3}$He/$^{4}$He $<$1). The period observed by {\sl ACE} on 2011 July 9 and later by {\sl STEREO-A} on 2011 July 16 was presumably associated with the small, compact AR 11246 located at the border of the coronal hole and showed very high $^{3}$He-enrichment and fluence, the highest so far detected on {\sl STEREO}. The source region produced EUV jets correlated with type III bursts and exhibited highly dynamic behavior. For example, when the active region temporarily disappeared the site continued with jet-like emissions. The characteristics of the events in this study are consistent with earlier suggestions that small flares have probably more favorable conditions for the $^{3}$He enrichments of SEPs than large flares \citep{rea88}, and that different size or morphology of the source active region with diverse flaring could produce a $^{3}$He-rich event \citep{kah87}. In spite of a dispersionless onset and absence of energetic electrons in the July 9 and 16 events the high cadence EUV images combined with the radio observations allowed us to determine the likely source brightening for the $^{3}$He emission. The approximate travel time of solar $^{3}$He ions with an energy 0.23-0.32~MeV\,nucleon$^{-1}$ is about 7 hours to {\sl ACE} along a nominal spiral with length of 1.2~AU. If the $^{3}$He was released during the eruption in AR 11246 at 16:30~UT on 2011 July 8 the ions with that energy should have been detected at the beginning of July 9, which is approximately consistent with the start of the event in the low energy time-intensity profile in Figure~\ref{fig3}b. Since the s/c was not connected to the coronal hole (and to the source AR) around the eruption time it missed the electrons and the first-arriving, higher-energy ions. This explains the lack of velocity dispersion in the ions received after the s/c enters the field lines connected to the coronal hole. If the release of the energetic ions for the 2011 July 16 period was associated for example with the jet and type III radio burst at 22:25~UT on July 15 then the lowest energy ions (0.18~MeV) shown in Figure~\ref{fig4}c should arrive at {\sl STEREO-A} around 06:30~UT on July 16 (i.e., three hours before the dispersionless event onset). If the ions were released during the brightening at 18:45~UT, which was associated with the prominent type III burst, the highest energy ions (1.49~MeV) should have only a $\sim$3~hr delay after the flare. This leads to the question of where these high energy solar ions, with enormous $^{3}$He enrichment, were residing in the heliosphere for about half of the day prior to the observed onset. The timing of the ion intensities in the 2011 July 16 period indicates that energetic $^{3}$He is closely associated with the CIR which is further supported by the pre- and post-event abundances. Such association may arise when a source active region is located at the periphery of the coronal hole with open field lines bending to the ecliptic. The high-speed solar wind emanating from the hole may create a corotating compression region in the heliosphere near $\sim$1 AU. Thus the energetic $^{3}$He injected from the solar source near the hole is guided by the open field to the CIR. The propagation of SEPs in the compressed field in CIRs has been modeled by \citet{koc03}. The authors employed the \citet{gia02} model of CIR acceleration on the solar wind speed gradients at 1~AU. Their simulations show that the magnetic enhancement associated with the CIR presents a kind of magnetic mirror away from the Sun where solar particles may be temporally trapped and re-accelerated. Indeed, our observations show a significant sunward component during the 2011 July 16 $^{3}$He-rich period indicating a reflecting boundary beyond the observer. \citet{koc08} have reported on multi-day $^{3}$He-rich periods that essentially all have dispersionless onset and were associated with compression in the solar wind but such convincing associations as in the July 16 period have not been presented. Those authors suggest that confinement of particles in CIRs is a significant factor of extended $^{3}$He-rich periods of SEPs. It is not clear if the weaker CIR seen during the 2011 July 9 $^{3}$He-rich period had a marked influence on the SEPs. The event commenced with a change of the magnetic connection and not with the start of the solar wind speed rise or at the CIR stream interface although the later was not fully developed. Presumably at larger radial distances where the corotating compression regions are found to be stronger, a closer association might exist. \citet{koc03} concluded that the effects of the compression may be important if the observed solar wind speed increase near Earth is more than 100~km\,s$^{-1}$ within a few hours. In the 2011 July 9 CIR, such a speed change is observed during a quite long $\sim$12~hr period. Notice that no significant concurrent CIR event was observed in association with the 2011 July 9 and July 16 compression corotating regions. That means at least for the strong July 16 compression that the bulk solar wind was not accelerated in the CIR. It is consistent with recent suggestions that CIR ions are accelerated out of the suprathermal ion pool \citep{mas12}. In the case of the July 16 CIR the suprathermal population was probably dominated by the $^{3}$He-rich SEPs which might undergo acceleration in the strong compression region. These small SEP events are often difficult to conclusively identify with source activity in active periods when there may be multiple candidates. In the July 9 event the s/c footpoint lies in a reasonable distance from both ARs 11246 and 11244. Note that the distance to AR 11244 was still within 3$\sigma$ ($\sim$48\degr) of the flare longitudes distribution associated with $^{3}$He-rich SEP events derived by \citet{rea99}. There are a few reasons to choose AR 11246 over AR 11244. The type III emissions at 11:37 and 16:25~UT preceding the July 9 event were probably associated with AR 11246 because of their temporal coincidence with the EUV brightenings in this AR; the brightenings, fainter at these times in AR 11244 were not new and likely continued from the previous activity. If the type III burst at 16:25~UT was associated with AR 11244 then why do we not detect escaping electrons, while being in a negative IMF sector and presumably connected to AR 11244? Another reason in favor of AR 11246 is the reversal of the IMF to the polarity at AR 11246 at the start of the July 9 event. Although AR 11244 continued with some activity on July 15-16, the type III bursts at 12:40, 14:15 and 18:45~UT on July 15 appear to be correlated with EUV brightenings in AR 11246. Note that some ambiguity may still remain in locating these type III bursts because we use lower (5 minute) cadence EUVI data. {\sl WIND}/WAVES observed weak and only low-frequency counterparts of these bursts. {\sl STEREO-B}/WAVES did not record any type III emission on July 15 in the same time period as {\sl STEREO-A}. This implies that the source was well hidden from the {\sl STEREO-B} view, which was not case of AR 11244. In addition, the AR 11244 location is less likely to be the source for the $^{3}$He-rich event on {\sl STEREO-A}. We have shown that the {\sl STEREO-A} footpoint was co-located with the coronal hole (containing AR 11246), while AR 11244 was quite far, 60\degr\ west of the footpoint. In conclusion, with the advantage of widely separated spacecraft carrying advanced imaging, radio and particle instrumentation, supported by the modeling of the coronal field, we have identified solar ARs which exhibited repeated energetic electron and $^{3}$He emissions for a relatively long interval of time, up to about a quarter of a solar rotation. This is significantly longer than single s/c observations which previously reported $^{3}$He injections from the same AR over a period of one day. Thus this study suggests that the conditions for the $^{3}$He acceleration in a single solar region may persist for a long time. Furthermore the observations in this paper show that a recurrent $^{3}$He-rich period can be preserved in the solar wind by confining SEPs in the CIR. The observations of a long-lived recurrent period presented in this study may provide new insights on the acceleration of $^{3}$He ions in the solar flares and their transport in the heliosphere. In particular, extended $^{3}$He-rich periods recently discovered with single s/c observations may be naturally explained by multiple injections from the same AR which may presumably last over days. | 14 | 3 | 1403.4856 |
1403 | 1403.6853_arXiv.txt | We present radial velocities, equivalent widths, model atmosphere parameters, and abundances or upper limits for 53~species of 48~elements derived from high resolution optical spectroscopy of 313~metal-poor stars. A majority of these stars were selected from the metal-poor candidates of the HK~Survey of Beers, Preston, and Shectman. We derive detailed abundances for 61\% of these stars for the first time. Spectra were obtained during a 10-year observing campaign using the Magellan Inamori Kyocera Echelle spectrograph on the Magellan Telescopes at Las Campanas Observatory, the Robert G.\ Tull Coud\'{e} Spectrograph on the Harlan J.\ Smith Telescope at McDonald Observatory, and the High Resolution Spectrograph on the Hobby-Eberly Telescope at McDonald Observatory. We perform a standard LTE abundance analysis using MARCS model atmospheres, and we apply line-by-line statistical corrections to minimize systematic abundance differences arising when different sets of lines are available for analysis. We identify several abundance correlations with effective temperature. A comparison with previous abundance analyses reveals significant differences in stellar parameters, which we investigate in detail. Our metallicities are, on average, lower by $\approx$~0.25~dex for red giants and $\approx$~0.04~dex for subgiants. Our sample contains 19~stars with [Fe/H]~$\leq -$3.5, 84~stars with [Fe/H]~$\leq -$3.0, and 210~stars with [Fe/H]~$\leq -$2.5. Detailed abundances are presented here or elsewhere for 91\% of the 209~stars with [Fe/H]~$\leq -$2.5 as estimated from medium resolution spectroscopy by Beers, Preston, and Shectman. We will discuss the interpretation of these abundances in subsequent papers. | \label{introduction} Observations of the high redshift Universe reveal carbon, magnesium, and other metals in the clouds of hydrogen that fueled the rising star formation rate density during the first several Gyr after the Big Bang (e.g., \citealt{sargent88,cooke11,matejek12}). Surely some of those gas clouds evolve into galaxies like today's Milky Way, whose stellar halos retain a chemical memory of those early epochs of star formation and metal production. If, however, one is interested in studying the early nucleosynthesis of less abundant metals, like holmium \citep{sneden96} or uranium \citep{cayrel01}, stars in the Milky Way are the only practical targets. For this and many other reasons, the importance of expanding the inventory of halo stars whose heavy metal abundances are known in great detail has long been recognized. \subsection{Previous Surveys of Metal-Poor Stars} Noteworthy in regard to the storyline of the present study are the surveys of \citet{bond70,bond80} and \citet{bidelman73}. The photographic plates for their objective-prism surveys were taken with the University of Michigan's 0.61~m Curtis Schmidt Telescope. This telescope, initially located near Ann Arbor, was relocated in 1966 to Cerro Tololo Inter-American Observatory (CTIO). Most of the observations for \citet{bond70} were made in Michigan, while \citeauthor{bidelman73} and \citet{bond80} made their observations at CTIO. Bond also used Str\"{o}mgren photometry to assign spectral types and luminosity classes to his candidates, and he measured radial velocities from followup coud\'{e} spectroscopy when possible. Many of the well-known and bright metal-poor stars in the Henry Draper (HD), Bonner Durchmusterung (BD), or C\'{o}rdoba Durchmusterung (CD) Catalogs were identified during this period by these surveys. Other metal-poor stars were found among the high proper motion stars in the Lowell Proper Motion Survey (\citealt{giclas71,giclas78}; these stars are identified with a ``G'' prefix before their catalog designation) and the New Luyten Two Tenths Catalog (NLTT, \citealt{luyten79}). \citet{ryan89} discusses the methods used to identify metal-poor stars in these surveys. \subsection{The HK Objective-Prism Survey} Our own work on the subject began with an objective-prism survey at the Curtis Schmidt (CS) Telescope at CTIO, initiated by G.\ Preston and S.\ Shectman, in 1978--1979. The key advance was the use of an interference filter to expose only the region around the stellar Ca~\textsc{ii}~H and K absorption lines at a spectral resolution of $\approx$~5~\AA\ (R~$\equiv \lambda/\Delta\lambda \sim$~800). This interference filter reduced crowding and sky fog and allowed longer exposure times (90~minutes) than had been practical previously. The ``HK'' Survey plates reach $B \approx$~15, several magnitudes fainter than the surveys of \citet{bond70,bond80} and \citet{bidelman73}. Visual inspection of the plates using a low-power binocular microscope yielded about 1800 metal-poor candidate stars on 72~plates. Broadband $UBV$ photometry and followup medium resolution (1~\AA; R~$\sim$~4,000) spectroscopy covering the 3700--4500~\AA\ wavelength range were obtained for 450~candidates. \citet{beers85} presented metallicity estimates for 134~metal-poor candidates with [Fe/H]~$< -$2.0. Using a revised metallicity calibration \citep{beers90}, Beers, Preston, \& Shectman (1992; hereafter \citeauthor{beers92}) published metallicity estimates, radial velocity measurements, and distances for 1044 dwarfs and giants with subsolar [Fe/H] from 135~unique fields covering 3375 square degrees (8\% of the sky). The relocation of Case Western Reserve University's Burrell Schmidt (BS) Telescope from Cleveland to Kitt Peak National Observatory (KNPO) in 1979 enabled a similar survey in the Northern Hemisphere. Followup medium resolution spectroscopy for these candidates was not yet available in 1992. \citet{beers13} describes the impressive worldwide network of 2--4~m class telescopes involved in the subsequent medium resolution spectroscopic followup of metal-poor candidates from both the Northern and Southern portions of the HK~Survey. Detailed abundances of metal-poor candidates from the Northern portion of the HK~Survey have been published elsewhere (e.g., \citealt{honda04a,honda04b}; \citealt{lai04,lai08}) and will not be considered in the sample presented here. An analysis of digital scans of the original HK~Survey plates was made by \citet{rhee01} and in subsequent unpublished work by J.\ Rhee, T.\ Beers, and coworkers (see also Section~3.3.1 of \citealt{beers05}). The goal of this ``HK-II'' Survey was to identify metal-poor red giant stars that may have been overlooked in the original visual scans of the plates due to the unavoidable temperature bias against cool stars. These candidates will not be considered in the sample presented here. \subsection{Subsequent Surveys} The Hamburg/ESO (HE) Survey \citep{wisotzki00} introduced quantitative methods to identify metal-poor candidates from digitized spectra \citep{christlieb08}. These techniques increase the effective yields of genuine metal-poor stars, especially among giants, when color information is included as part of the selection criteria. This survey also reaches several magnitudes deeper than the HK~Survey. Recent surveys have built on these quantitative techniques to identify even greater numbers of candidate metal-poor stars from low resolution spectroscopy, including the Radial Velocity Experiment (RAVE; see \citealt{fulbright10}), the Sloan Digital Sky Survey (SDSS), and the Sloan Extensions for Galactic Understanding and Exploration (SEGUE-1 and SEGUE-2; \citealt{yanny09,rockosi12}). Many metal-poor candidates are also expected to be found among ongoing surveys by the Large sky Area Multi-Object fiber Spectroscopic Telescope (LAMOST; e.g., \citealt{deng12}); and the SkyMapper Telescope \citep{keller12}. \subsection{Detailed Abundance Followup of Metal-Poor Candidates} Bright stars identified by the surveys of \citet{bond70,bond80} and \citet{bidelman73} have been analyzed in great chemical detail by numerous investigators, including \citet{luck81,luck85}, \citet{hartmann88}, \citet{gilroy88}, \citet{gratton88,gratton91,gratton94}, \citet{magain89}, \citet{peterson90}, \citet{zhao90}, and \citet{johnson02}. \citet{ryan91b}, \citet{fulbright00,fulbright02}, \citet{stephens02}, and \citet{ishigaki10} studied the compositions of stars selected from the early objective prism and proper motion surveys. Detailed abundance studies of large samples of stars from the HE Survey have been conducted by \citet{carretta02}, \citet{cohen02,cohen04,cohen08}, \citet{barklem05}, \citet{aoki07}, \citet{hollek11}, \citet{norris13}, and \citet{yong13}. Candidates from the HK-II Survey have been observed as part of the Chemical Abundances of Stars in the Halo (CASH) project at the University of Texas \citep{frebel08b,roederer08b}. Detailed chemical followup of large numbers of stars from the SDSS and SEGUE have been performed by \citet{aoki08,aoki13}, \citet{caffau11}, and \citet{bonifacio12}. Over the last several decades, high resolution optical spectroscopic followup of candidates from \citeauthor{beers92} has confirmed hundreds of them as genuine metal-poor stars. Chemical abundances of handfuls of stars from Beers et al.\ (\citeyear{beers85}) were presented by \citet{molaro90a}, \citet{molaro90b}, \citet{norris93}, and \citet{primas94}. \citet{mcwilliam95a,mcwilliam95b}, \citet{norris96}, and \citet{ryan96} were the first to analyze larger samples of stars (34~stars between them) from \citeauthor{beers92}.~ Since then, the number of detailed abundance studies conducted on candidates from \citeauthor{beers92} has grown tremendously, and there are far too many excellent ones to list here individually. There have been several dedicated observing campaigns to obtain high resolution spectroscopy of substantial numbers of stars (typically $\approx$~10--30), including analyses by \citet{aoki05,aoki07}, \citet{honda04a,honda04b}, \citet{lai08}, the ``First Stars'' team \citep{cayrel04,spite05,francois07,bonifacio09}, and a reanalysis of the published values from many of these studies by \citet{yong13}. The detailed chemical analysis performed by \citet{mcwilliam95a,mcwilliam95b} launched our efforts to use these stars as probes of the earliest epoch of metal enrichment in the Galaxy. Our subsequent abundance studies based on high resolution spectroscopy of metal-poor candidates from the HK~Survey have examined carbon rich metal-poor stars \citep{preston01,sneden03b,roederer14}, individual stars of interest \citep{sneden94,ivans05,preston06b,thompson08}, stars on the horizontal branch \citep{preston06}, and stars with kinematics indicative of a cold stellar stream \citep{roederer10}. In this paper we present abundance results for 313~stars, including 217~stars from the HK~Survey, using high resolution spectroscopy obtained from 2003--2013 at the Magellan Telescopes at Las Campanas Observatory. As of July, 2013, detailed abundances for 91\% (191/209) of the stars with estimated [Fe/H]~$\leq -$2.5 in \citeauthor{beers92} are presented here or have been published elsewhere previously. Sixty-one percent (132/217) of the stars from the HK~Survey presented in this work are analyzed in such a manner for the first time. Abundances in the other 85~stars have been examined previously by the studies above or others named below. A limited selection of stars from the BD, CD, G, HD, and HE catalogs are also (re)analyzed in the present study. | 14 | 3 | 1403.6853 |
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1403 | 1403.1368_arXiv.txt | We present a method for distance calibration without using standard fitting procedures. Instead we use random resampling to reconstruct the probability density function (PDF) of calibration data points in the fitting plane. The resulting PDF is then used to estimate distance-related properties. The method is applied to samples of radio surface brightness to diameter ($\Sigma-D$) data for the Galactic supernova remnants (SNRs) and planetary nebulae (PNe), and period-luminosity ($PL$) data for the Large Magellanic Cloud (LMC) fundamental mode classical Cepheids. We argue that resulting density maps can provide more accurate and more reliable calibrations than those obtained by standard linear fitting procedures. For the selected sample of the Galactic SNRs, the presented PDF method of distance calibration results in a smaller average distance fractional error of up to $\approx16$ percentage points. Similarly, the fractional error is smaller for up to $\approx8$ and $\approx0.5$ percentage points, for the samples of Galactic PNe and LMC Cepheids, respectively. In addition, we provide a PDF-based calibration data for each of the samples. | Scaling relations are widely used, sometimes as the only option, to determine the relevant properties of astrophysical objects. One particularly important property in astrophysical studies is the distance to a particular object. Direct measurement of distances is often not possible, and the only way to infer the distance is from a scaling relation. The calibration of these relations is very important and extensive scientific efforts have been made to assess the quality of the calibration data samples and applied calibration procedures. Most commonly, sample of calibrators is fitted with some analytical functional dependence, where one of the data variables is dependent on one or more remaining data variables. In cases where all data variables have significant uncertainties and cannot be resolved to dependent/independent ones, fitting procedures are adjusted accordingly. If there are $N$ calibrators in a particular calibrating sample, with $n$ coordinates per data point, then there are $ n \times N $ numbers of information in that particular calibrating sample. When fitting, all of this information is projected into the linear fit parameters. The initial information contained in the calibration sample is thus reduced and averaged out. While this might not be problematic for the samples with strong functional dependence, in the case of loose correlations (often used in astrophysics) it might cause a significant difference, depending on the type of functional dependence, type of offsets from the best fit line and other assumptions of the applied fitting procedure \citep{Isobe_etal_1990ApJ, urosevic_etal_10, pavlovicetal13}. Modelling kernels based on analytical methods often require that the underlying phenomena appear smooth and predictable. This, however, need not be the case, especially when simplifications are invoked to describe the complexity of the systems with a small number of parameters. This is usually done in studies of astrophysical objects. Evolution of such objects, despite their complex macroscopic appearance, reflected in diversity of intrinsic and environmental parameters, is often described with only two parameters, as is the case with calibration relations considered in this paper. The intrinsic complexity of the nature itself leads to events that cannot be predicted using a simplified analytic approach \citep{Taleb2007}. Consequently, the shape of the data sample probability density function (PDF) might not follow the direction of the same data sample best fit line and can deviate from that line in Gaussian manner, even in cases of complete and well-studied samples. Our approach relies on numerical calculation of PDF of calibrating data rather than on applying a fitting procedure to them. This increases the likelihood that the information contained in the calibration sample is preserved and ensures greater consistency and more accurate calibrations. The standard procedure to estimate PDF from data samples is to bin the data and make histograms. The problem with this approach is that the reconstructed PDFs can heavily depend on the bin size, especially when incomplete calibrating samples are considered { \citep[for more information on assessing the binning problem, we refer the reader to][where a sophisticated approach with Bayesian Blocks method is used]{Scargle_et_al_2013ApJ_764_167S}}. Another possible approach is to reconstruct cumulative distribution function (CDF) and then reconstruct the PDF \citep{BergHarris08}. However, CDF reconstruction requires data sorting and can only be performed on unidimensional data. Our approach requires no histograms and no CDF reconstruction. We calculate PDF using Monte Carlo resampling of the calibration (original) sample. Coordinates of the data points in resampled samples are translated to account for the difference between centroid coordinates of the original and resampled sample. Our algorithm stems from the basic principles of bootstrap statistics \citep[random resampling,][]{EfronTibshirani93} and principal component analysis \citep[standardization of data coordinates using centroids and calculation of highest data variability direction,][]{pearson1901, Jolliffe_PCA_2002}. Albeit simple, the algorithm is computationally intensive. However, in present times of ubiquitous computing resources it can be performed with sufficient accuracy on a standard office computer and it can yield smooth PDFs that resemble the distribution of calibration sample data points. Even small samples (of $\sim 10 $ data points) can give smooth density maps of high resolution (see Equation \ref{n_resampl} and Figure \ref{pdf_snrs_pne}). This can be very significant for calibrating relations where scarce samples are a rule rather than exception. We developed our algorithm for the purpose of calibrating bidimensional data samples and our analysis and algorithm presentation will be constrained to two dimensions. The simplicity of our approach makes a multidimensional data application a quite straightforward extension (see Section 2). We apply our analysis to the radio surface brightness to diameter { ($\Sigma-D$)} relation for supernova remnants (SNRs) and planetary nebulae (PNe), and also to the period-luminosity ($PL$) relation for classical Cepheids. The majority of the work done on the $\Sigma-D$ relation was in order to produce reliable calibrations that can be used to calculate distances. Samples often suffer from significant scatter and if the information that they contain is described with the parameters of the best fit line, it can result in inaccurate distance estimates. On the other hand, the $PL$ relation gives more consistent distances, with the order of magnitude smaller average fractional error than the above mentioned $\Sigma-D$ relations. The average fractional error for distance is calculated as: \begin{equation} \bar{f}=\frac{1}{N} \sum_{i=1}^N \left| \frac{d_\mathrm{i}-d_\mathrm{i}^\mathrm{s}}{d_\mathrm{i}} \right|, \end{equation} where $N$ is the number of data points in the calibrating sample, $d_\mathrm{i}$ is the measured distance to the object represented with $\mathrm{i}^\mathrm{th}$ data point and $d_\mathrm{i}^\mathrm{s}$ is a statistical distance to that object determined either from the best fit line to the calibrating data set or in some other way such as the PDF-based method presented in this paper. \subsection{The $\Sigma-D$ relations for supernova remnants and planetary nebulae} The relation between the radio surface brightness and diameter is usually given as: \begin{equation} \Sigma = AD^{-\beta}, \label{SigD} \end{equation} where $A$ and $\beta$ are parameters. This is the standard form that follows from theoretical work \citep[first derived for supernova remnants, ][]{shklovskii60a} and is readily used for calibration. Calibration is performed by linearising the above equation and applying some of the standard fitting techniques \citep{shklovskii60b}. Also, from theoretical considerations it is expected that $A$ and $\beta$ have different values in different stages of SNR evolution and this can also interfere with calibration precision when modelling SNR evolution with only one evolutionary trajectory (the best fit line), which is usually done. Once calibrated, the relation can be used to determine the distance to a particular SNR by measuring its flux density $S = \Sigma / \Omega$ and angular diameter $\theta = \sqrt{4\Omega/\pi}$. After calculating $\Sigma$, the corresponding value for $D$ follows from Equation \ref{SigD} and distance can be calculated as $d=\theta D$. The $\Sigma-D$ relation for SNRs has more than five decades' long history. In addition to further theoretical development \citep[i.e.,][]{DuricSeaquist86}, extensive work was done on calibrating the relation for distance determination \citep[for some calibrations see,][]{urosevic_etal05,CaseBattacharya98,allakhverdievetal86}. The $\Sigma-D$ relation for planetary nebulae in the form of Equation \ref{SigD} was theoretically derived and empirically assessed by \citet{urosevic_etal07,urosevic_etal09b}. The above papers use standard fitting procedures based on vertical offsets and are mostly not concerned with further development of fitting procedures. \citet{pavlovicetal13} argued that applying different types of fitting offsets can result in different parameters of the $\Sigma-D$ relation for SNRs, and that orthogonal offsets are more reliable and stable over other types of offsets. Similar analysis, but to a lesser extent, was performed on a PNe sample from \citet{stanghellini_etal08} in \citet{VukoticUrosevic12}. Although these analyses argue in favour of orthogonal offsets calibrations, the dependence of calibration parameters on the type of selected fitting offsets introduces further ambiguities in the efforts towards reliable calibrations. Also, poor quality of the calibrating samples often results in statistically unacceptable fits. The calibration algorithm proposed in this paper is not using fitting procedures and there are no assumptions on the type of functional dependence in calibrating relation. This makes the resulting calibration more consistent, with no loss of information, because the initial information contained in the data points coordinates is not reduced to the parameters of the best fit line. Here we apply our algorithm for data density distribution calculation to the sample of 60 Galactic SNRs from \citet{pavlovicetal13} and to the 39 Galactic PNe with reliable distances from \citet{stanghellini_etal08}. \subsection{The $PL$ relation for Cepheids} \label{intro_ceph} In the case of Cepheids, pulsating variable stars, historically the period-luminosity relation \citep{Leavitt_1912} was of crucial importance for determining distances. The $PL$ relation is also a starting point of the distance ladder \citep[e.g.][]{Rowan_1985}. Using Cepheids to estimate distances to other galaxies is one of the starting points in measuring the Hubble constant ($H_0$). Small $PL$ calibration inaccuracies can propagate to significantly larger discrepancies in estimates of the universe expansion rate. As the body of data increased, it became obvious that there are some problems with $PL$ relation, some of which remain unsolved until this day \citep{Sandage_tammann2006}. The problems of the $PL$ relation are connected to the problem of reddening in different direction of eyesight, investigation of the metallicity dependence, phase dependency of the relation, universality of the $PL$ relation, or study of the non-linearity of the $PL$ function of various samples \citep[e.g.][]{Garcia_Varela_2013,Kanbur_2010,Koen_2007}. The $PL$ relations in mid infra-red are somewhat less problematic than their visual counterparts, but still, relations at $3.6~{\mu}m$ and $4.6~{\mu}m$ for samples of Large Magellanic Cloud (LMC) Cepheids \citep[]{Freedman_2008,Ngeow_2008,scowcroft_etal_11} leave open questions about their metallicity dependency. All of these issues directly affect the estimation of the $H_0$. Studies on the $PL$ relation usually present luminosity in the form of an absolute magnitude $M$ and use $M-\log{P}$ form of the data for plotting and fitting. Once calibrated, the $M-\log{P}$ form can be used as a distance estimation tool. Similarly to the $\Sigma-D$ relation, if the period of a pulsating star is measured, a corresponding $M$ value that is derived from the parameters of the calibration (the line fitted to the calibration sample) can be used to estimate distance $d$ to an object with a measured value of apparent magnitude ($m_\mathrm{o}$): \begin{equation} d = 10^{\left( 1.0 + \frac{m_\mathrm{o}-M}{5.0}\right)}. \label{distance_PL} \end{equation} We selected the LMC fundamental mode Cepheid samples in I and V band from OGLE project \citep{Soszynski_etal2008AcA}, that was corrected for extinction by \citet[][and references therein]{Ngeow_etal_ApJ09}. Compared to the considered samples of Galactic SNRs and PNe, these samples have a significantly larger number of data points { (better plotting plane coverage)} and should yield more accurate distance calibrations. { Also, better accuracy is evident from the smaller scatter in the Cepheid samples relative to the selected axis range, than in the case of SNR and PN samples (scatter of the PDF signal from the best fit line in Figures \ref{pdf_snrs_pne} and \ref{ceph_pdf_I_V}). Initial conditions and host environments of SNRs and PNe are by far more diverse than for Cepheids and consequently less accurate when described with a single linear relation.} This paper is organized as follows. The next section describes the implementation of our PDF-based algorithm and its features. In Section \ref{analysis} we present the resulting PDF for selected samples of SNRs, PNe and Cepheids, respectively, and discuss the results. Also, in Table 1 we give the average fractional errors of our calibrations and compare them with previous calibrations while our PDF-based calibrations themselves are given in Tables 2 and 3. Summary of the results and conclusions of this paper are presented in the last section. | We presented a calibration method which relies on density distribution of data points rather than fitting procedures. The resulting calibrations are more robust and accurate and require no assumptions on functional dependence such as fitting-based calibrations. Our algorithm for generating data sample PDF is based on calculating centroid offsets between Monte Carlo resampled and the actual calibrating data sample. As such, it requires no binning such as in histogram-based approaches. The method is applied to distance related scaling relations, the $\Sigma-D$ relation for SNRs and PNe and the $PL$ relation for Cepheids. The selected samples of Galactic SNRs and PNe have a much larger scatter than selected samples of fundamental mode LMC Cepheids in I and V band. This is due to the fact that Galactic distance mapping is tenuous compared to the case where all objects in the sample reside in external galaxy and are approximately at the same distance. Compared to the best calibrating fit lines for the selected SNRs and PNe, our method gives up to $\approx 16$ and $\approx 8$ percentage points, respectively, smaller average fractional error for distance estimates. Even in the case of LMC Cepheids, a mere fractional error reduction of up to $\approx 0.5$ percentage points can be a significant improvement in building up a Cepheid-based distance ladder. Apart from improvement in accuracy, the PDF-based calibrating method presented in this paper gives much more information about the data sample than standard fitting techniques. In the proposed method, the information contained in the calibrating samples is preserved rather than averaged out and condensed into the parameters of the best fit line. In the case where data samples are reliable and complete, this preserved information can be used to give more insight into the evolution of the examined objects. This could be a viable tool for quantifying the dependence of the Galactic SNRs evolution on the ambient medium density or some other relevant feature and the same holds for Galactic PNe. In the case of LMC Cepheids, small deviations in a sample PDF can be used to trace the fine details of the nature of pulsations and improvement of the distance ladder. The main purpose of this paper was to present a more accurate method for statistical distance determination. The more informative nature of the method opens new vistas in giving clues on the evolution of different objects through quantification of sample PDF features. We leave further development of the method in this direction for future work. \begin{table*} \label{calibtable} \centering \begin{minipage}{170mm} \caption{Calibrating values for distance determination for the Galactic $\Sigma-D$ relations (SNRs and PNe) and the $PL$ relation (in the form of apparent magnitude $m$ vs. period $P$) for fundamental mode Cepheids in I and V band. The units of the $\log$ values and coordinate range are the same as in Figures \ref{pdf_snrs_pne} and \ref{ceph_pdf_I_V}. The grid size is $10^2 \times 10^2$.} \begin{tabular}{ccccccccccccc@{\hspace*{-0.1mm}}ccc} \hline \multicolumn{4}{c}{--------------- SNRs ---------------}&\multicolumn{4}{c}{--------------- PNe ---------------}&\multicolumn{8}{c}{----------------------------- Cepheids -----------------------------}\\ $\log{\Sigma}$ &\multicolumn{3}{c}{$D[pc]$}&$\log{\Sigma}$&\multicolumn{3}{c}{$D[pc]$}&$\log{P}$&\multicolumn{3}{c}{$m_{\mathrm{I}}$}&&\multicolumn{3}{c}{$m_{\mathrm{V}}$}\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} &mode&mean&med.&&mode&mean&med.&&mode&mean&med.&&mode&mean&med.\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} \hline -23.42& -- & -- & -- & -24.90& -- & -- & -- & -0.08& -- & -- & -- & & -- & -- & -- \\ -23.34& -- & -- & -- & -24.80& -- & -- & -- & -0.06& -- & -- & -- & & -- & -- & -- \\ -23.26& -- & -- & -- & -24.70& -- & -- & -- & -0.05& -- & -- & -- & & -- & -- & -- \\ -23.18& -- & -- & -- & -24.60& -- & -- & -- & -0.03& -- & -- & -- & & -- & -- & -- \\ -23.10& -- & -- & -- & -24.50& -- & -- & -- & -0.01& 16.69& 16.69& 16.69& & 17.18& 17.18& 17.18\\ -23.02& -- & -- & -- & -24.40& -- & -- & -- & 0.01& 16.62& 16.62& 16.69& & 17.11& 17.11& 17.18\\ -22.94& -- & -- & -- & -24.30& -- & -- & -- & 0.03& 16.48& 16.55& 16.48& & 16.97& 17.04& 16.97\\ -22.86& 44.67& 44.67& 44.67& -24.20& -- & -- & -- & 0.04& 16.55& 16.48& 16.55& & 17.11& 17.04& 17.11\\ -22.78& 44.67& 41.69& 44.67& -24.10& -- & -- & -- & 0.06& 16.41& 16.41& 16.48& & 16.83& 16.83& 16.90\\ -22.70& 41.69& 38.90& 41.69& -24.00& 7.24& 7.94& 7.94& 0.08& 16.34& 16.34& 16.34& & 16.83& 16.83& 16.83\\ -22.62& 38.90& 38.90& 38.90& -23.90& 7.24& 7.24& 7.94& 0.10& 16.34& 16.41& 16.34& & 16.83& 16.90& 16.83\\ -22.54& 36.31& 36.31& 38.90& -23.80& 7.24& 7.24& 7.24& 0.12& 16.27& 16.41& 16.34& & 16.83& 16.83& 16.83\\ -22.46& 36.31& 36.31& 36.31& -23.70& 6.61& 6.61& 6.61& 0.13& 16.27& 16.27& 16.27& & 16.69& 16.69& 16.69\\ -22.38& 33.88& 33.88& 36.31& -23.60& 6.03& 6.03& 6.03& 0.15& 16.13& 16.13& 16.20& & 16.69& 16.69& 16.69\\ -22.30& 33.88& 33.88& 33.88& -23.50& 5.50& 5.50& 6.03& 0.17& 16.27& 16.20& 16.13& & 16.62& 16.69& 16.62\\ -22.22& 31.62& 31.62& 31.62& -23.40& 5.50& 5.50& 5.50& 0.19& 15.99& 16.13& 16.06& & 16.41& 16.62& 16.62\\ -22.14& 31.62& 31.62& 31.62& -23.30& 5.01& 5.01& 5.01& 0.21& 16.06& 15.99& 16.06& & 16.55& 16.55& 16.55\\ -22.06& 29.51& 29.51& 29.51& -23.20& 4.57& 4.57& 4.57& 0.22& 16.06& 16.06& 16.06& & 16.48& 16.62& 16.62\\ -21.98& 29.51& 31.62& 29.51& -23.10& 4.17& 4.17& 4.57& 0.24& 15.99& 15.99& 15.99& & 16.48& 16.48& 16.48\\ -21.90& 15.85& 38.90& 27.54& -23.00& 4.17& 4.17& 4.17& 0.26& 15.85& 15.92& 15.92& & 16.34& 16.41& 16.48\\ -21.82& 14.79& 44.67& 36.31& -22.90& 3.80& 3.80& 3.80& 0.28& 15.92& 15.85& 15.92& & 16.48& 16.41& 16.48\\ -21.74& 14.79& 51.29& 67.61& -22.80& 3.47& 3.47& 3.47& 0.30& 15.64& 15.78& 15.78& & 16.20& 16.34& 16.34\\ -21.66&165.96& 51.29& 63.10& -22.70& 3.16& 3.47& 3.47& 0.31& 15.78& 15.71& 15.71& & 16.20& 16.20& 16.20\\ -21.58& 33.88& 51.29& 58.88& -22.60& 3.16& 3.16& 3.47& 0.33& 15.71& 15.64& 15.71& & 16.27& 16.20& 16.27\\ -21.50& 31.62& 47.86& 54.95& -22.50& 3.16& 3.16& 3.16& 0.35& 15.78& 15.64& 15.64& & 16.41& 16.20& 16.20\\ -21.42& 31.62& 47.86& 54.95& -22.40& 2.88& 2.88& 3.16& 0.37& 15.29& 15.43& 15.50& & 15.71& 15.99& 16.06\\ -21.34& 29.51& 47.86& 54.95& -22.30& 2.88& 2.63& 2.88& 0.39& 15.50& 15.43& 15.50& & 15.99& 15.99& 15.99\\ -21.26& 54.95& 44.67& 54.95& -22.20& 2.63& 2.19& 2.63& 0.40& 15.43& 15.43& 15.43& & 16.06& 15.99& 15.99\\ -21.18& 54.95& 38.90& 38.90& -22.10& 2.40& 1.82& 2.40& 0.42& 15.36& 15.36& 15.36& & 15.92& 15.92& 15.99\\ -21.10& 31.62& 33.88& 33.88& -22.00& 1.05& 1.51& 1.51& 0.44& 15.29& 15.29& 15.36& & 15.85& 15.85& 15.92\\ -21.02& 29.51& 29.51& 31.62& -21.90& 0.95& 1.38& 1.26& 0.46& 15.22& 15.29& 15.29& & 15.78& 15.85& 15.85\\ -20.94& 29.51& 29.51& 29.51& -21.80& 0.87& 1.26& 1.15& 0.48& 15.22& 15.22& 15.22& & 15.78& 15.78& 15.78\\ -20.86& 47.86& 29.51& 29.51& -21.70& 0.87& 1.15& 1.05& 0.49& 15.15& 15.15& 15.15& & 15.71& 15.78& 15.78\\ -20.78& 44.67& 29.51& 33.88& -21.60& 0.79& 1.05& 0.95& 0.51& 15.08& 15.15& 15.15& & 15.71& 15.71& 15.71\\ -20.70& 41.69& 29.51& 33.88& -21.50& 0.72& 0.95& 0.87& 0.53& 15.08& 15.08& 15.08& & 15.64& 15.71& 15.71\\ -20.62& 41.69& 29.51& 31.62& -21.40& 0.66& 0.95& 0.79& 0.55& 15.08& 15.01& 15.08& & 15.64& 15.64& 15.64\\ -20.54& 38.90& 27.54& 29.51& -21.30& 0.60& 0.87& 0.79& 0.57& 15.01& 14.94& 15.01& & 15.50& 15.57& 15.57\\ -20.46& 25.70& 27.54& 27.54& -21.20& 0.60& 0.87& 0.95& 0.58& 14.87& 14.87& 14.94& & 15.57& 15.50& 15.57\\ -20.38& 25.70& 27.54& 27.54& -21.10& 0.95& 0.87& 0.95& 0.60& 14.87& 14.87& 14.87& & 15.29& 15.43& 15.50\\ -20.30& 23.99& 27.54& 25.70& -21.00& 0.87& 0.79& 0.87& 0.62& 14.87& 14.80& 14.87& & 15.57& 15.43& 15.50\\ -20.22& 22.39& 27.54& 25.70& -20.90& 0.79& 0.79& 0.87& 0.64& 14.66& 14.73& 14.80& & 15.22& 15.36& 15.36\\ -20.14& 33.88& 29.51& 31.62& -20.80& 0.79& 0.79& 0.79& 0.66& 14.73& 14.73& 14.73& & 15.36& 15.36& 15.36\\ -20.06& 33.88& 29.51& 31.62& -20.70& 0.72& 0.72& 0.72& 0.67& 14.66& 14.66& 14.66& & 15.22& 15.29& 15.29\\ -19.98& 31.62& 27.54& 29.51& -20.60& 0.66& 0.72& 0.79& 0.69& 14.59& 14.59& 14.66& & 15.29& 15.29& 15.29\\ -19.90& 27.54& 27.54& 27.54& -20.50& 0.79& 0.72& 0.79& 0.71& 14.38& 14.52& 14.52& & 14.94& 15.08& 15.15\\ -19.82& 27.54& 27.54& 27.54& -20.40& 0.72& 0.72& 0.72& 0.73& 14.45& 14.45& 14.52& & 15.15& 15.08& 15.15\\ -19.74& 25.70& 27.54& 27.54& -20.30& 0.66& 0.66& 0.66& 0.75& 14.45& 14.38& 14.45& & 15.01& 15.01& 15.01\\ -19.66& 25.70& 27.54& 27.54& -20.20& 0.60& 0.60& 0.60& 0.76& 14.45& 14.38& 14.45& & 15.22& 15.01& 15.08\\ -19.58& 31.62& 27.54& 29.51& -20.10& 0.60& 0.55& 0.55& 0.78& 14.24& 14.24& 14.24& & 14.87& 14.87& 14.87\\ -19.50& 29.51& 23.99& 27.54& -20.00& 0.38& 0.55& 0.50& 0.80& 14.17& 14.24& 14.24& & 14.87& 14.87& 14.87\\ -19.42& 9.12& 20.89& 23.99& -19.90& 0.35& 0.50& 0.38& 0.82& 14.17& 14.24& 14.24& & 15.01& 14.87& 14.94\\ -19.34& 23.99& 18.20& 22.39& -19.80& 0.24& 0.46& 0.35& 0.84& 14.17& 14.10& 14.17& & 14.80& 14.73& 14.80\\ -19.26& 9.77& 16.98& 20.89& -19.70& 0.24& 0.38& 0.35& 0.85& 14.03& 14.03& 14.10& & 14.80& 14.66& 14.73\\ -19.18& 9.77& 14.79& 15.85& -19.60& 0.29& 0.35& 0.32& 0.87& 14.03& 13.96& 14.03& & 14.24& 14.59& 14.59\\ -19.10& 9.77& 12.88& 14.79& -19.50& 0.26& 0.32& 0.32& 0.89& 13.96& 13.96& 13.96& & 14.52& 14.59& 14.59\\ \hline \end{tabular} \end{minipage} \end{table*} \addtocounter{table}{-1} \begin{table*} \centering \begin{minipage}{170mm} \caption{-- {\it continued}} \begin{tabular}{ccccccccccccc@{\hspace*{-0.1mm}}ccc} \hline \multicolumn{4}{c}{--------------- SNRs ---------------}&\multicolumn{4}{c}{--------------- PNe ---------------}&\multicolumn{8}{c}{----------------------------- Cepheids -----------------------------}\\ $\log{\Sigma}$ &\multicolumn{3}{c}{$D[pc]$}&$\log{\Sigma}$&\multicolumn{3}{c}{$D[pc]$}&$\log{P}$&\multicolumn{3}{c}{$m_{\mathrm{I}}$}&&\multicolumn{3}{c}{$m_{\mathrm{V}}$}\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} &mode&mean&med.&&mode&mean&med.&&mode&mean&med.&&mode&mean&med.\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} \hline -19.02& 14.79& 9.77& 10.47& -19.40& 0.42& 0.32& 0.32& 0.91& 13.96& 13.89& 13.96& & 14.59& 14.52& 14.59\\ -18.94& 9.77& 8.51& 9.77& -19.30& 0.38& 0.29& 0.32& 0.93& 13.68& 13.82& 13.89& & 14.87& 14.52& 14.52\\ -18.86& 4.27& 7.41& 8.51& -19.20& 0.35& 0.29& 0.32& 0.94& 13.68& 13.75& 13.75& & 14.38& 14.38& 14.38\\ -18.78& 5.25& 6.46& 5.62& -19.10& 0.35& 0.26& 0.29& 0.96& 13.75& 13.82& 13.75& & 14.38& 14.45& 14.45\\ -18.70& 5.25& 6.46& 5.62& -19.00& 0.32& 0.22& 0.24& 0.98& 13.82& 13.75& 13.82& & 14.59& 14.38& 14.52\\ -18.62& 10.47& 6.92& 6.03& -18.90& 0.17& 0.18& 0.20& 1.00& 13.47& 13.68& 13.75& & 14.17& 14.38& 14.38\\ -18.54& 10.47& 7.94& 9.77& -18.80& 0.15& 0.15& 0.17& 1.02& 13.54& 13.61& 13.61& & 14.24& 14.31& 14.31\\ -18.46& 10.47& 8.51& 10.47& -18.70& 0.15& 0.13& 0.15& 1.03& 13.68& 13.61& 13.68& & 14.45& 14.24& 14.45\\ -18.38& 9.77& 9.77& 10.47& -18.60& 0.14& 0.13& 0.14& 1.05& 13.33& 13.47& 13.47& & 14.38& 14.10& 14.17\\ -18.30& 9.77& 10.47& 11.22& -18.50& 0.14& 0.13& 0.14& 1.07& 13.40& 13.40& 13.40& & 14.10& 14.10& 14.17\\ -18.22& 12.02& 10.47& 11.22& -18.40& 0.13& 0.14& 0.14& 1.09& 13.47& 13.47& 13.47& & 14.24& 14.24& 14.24\\ -18.14& 11.22& 10.47& 11.22& -18.30& 0.11& 0.14& 0.14& 1.11& 13.47& 13.40& 13.47& & 14.17& 14.10& 14.17\\ -18.06& 11.22& 10.47& 11.22& -18.20& 0.11& 0.14& 0.14& 1.12& 13.26& 13.26& 13.33& & 13.96& 14.03& 14.03\\ -17.98& 10.47& 10.47& 11.22& -18.10& 0.08& 0.14& 0.13& 1.14& 13.33& 13.26& 13.33& & 14.10& 14.03& 14.10\\ -17.90& 10.47& 10.47& 10.47& -18.00& 0.07& 0.13& 0.10& 1.16& 13.19& 13.19& 13.19& & 13.89& 14.03& 13.96\\ -17.82& 10.47& 10.47& 10.47& -17.90& 0.07& 0.11& 0.09& 1.18& 13.26& 13.19& 13.12& & 14.03& 13.89& 13.89\\ -17.74& 9.77& 9.77& 10.47& -17.80& 0.06& 0.10& 0.07& 1.20& 13.26& 13.19& 13.26& & 13.89& 13.96& 13.89\\ -17.66& 9.12& 9.12& 9.12& -17.70& 0.06& 0.08& 0.07& 1.21& 13.26& 13.19& 13.26& & 13.47& 13.89& 14.03\\ -17.58& 9.12& 9.12& 9.12& -17.60& 0.05& 0.07& 0.06& 1.23& 13.19& 13.12& 13.19& & 13.96& 13.82& 13.96\\ -17.50& -- & -- & -- & -17.50& 0.05& 0.06& 0.05& 1.25& 13.05& 12.98& 13.05& & 13.82& 13.82& 13.82\\ -17.42& 6.46& 6.92& 6.92& -17.40& 0.10& 0.05& 0.05& 1.27& 12.84& 12.91& 12.84& & 13.61& 13.68& 13.68\\ -17.34& 6.46& 6.46& 6.46& -17.30& 0.09& 0.05& 0.08& 1.29& 12.98& 12.84& 12.98& & 13.89& 13.61& 13.68\\ -17.26& 6.03& 6.03& 6.03& -17.20& 0.08& 0.05& 0.07& 1.30& 12.56& 12.56& 12.56& & 13.89& 13.61& 13.47\\ -17.18& 6.03& 6.03& 6.03& -17.10& 0.07& 0.05& 0.07& 1.32& 12.84& 12.70& 12.84& & 13.54& 13.47& 13.54\\ -17.10& 5.62& 5.62& 5.62& -17.00& 0.07& 0.05& 0.07& 1.34& 12.70& 12.77& 12.77& & 13.54& 13.40& 13.54\\ -17.02& 5.62& 5.62& 5.62& -16.90& 0.06& 0.05& 0.07& 1.36& 12.77& 12.63& 12.70& & 13.26& 13.40& 13.47\\ -16.94& 5.25& 5.25& 5.25& -16.80& 0.05& 0.05& 0.06& 1.38& 12.63& 12.63& 12.63& & 13.40& 13.33& 13.40\\ -16.86& 4.90& 4.90& 5.25& -16.70& 0.05& 0.05& 0.05& 1.39& 12.49& 12.49& 12.49& & 13.33& 13.33& 13.33\\ -16.78& 4.90& 4.90& 4.90& -16.60& 0.05& 0.05& 0.05& 1.41& 12.42& 12.42& 12.42& & 13.47& 13.33& 13.47\\ -16.70& 4.57& 4.57& 4.90& -16.50& 0.05& 0.05& 0.05& 1.43& 12.42& 12.49& 12.42& & 12.77& 13.05& 12.84\\ -16.62& 4.57& 4.57& 4.57& -16.40& 0.04& 0.04& 0.05& 1.45& 12.49& 12.49& 12.49& & 13.33& 13.12& 13.12\\ -16.54& 4.57& 4.57& 4.57& -16.30& 0.04& 0.04& 0.04& 1.47& 12.49& 12.42& 12.49& & 13.05& 13.05& 13.05\\ -16.46& 4.27& 4.27& 4.57& -16.20& 0.04& 0.04& 0.04& 1.48& 12.28& 12.35& 12.28& & 13.05& 12.84& 13.05\\ -16.38& 4.27& 4.27& 4.27& -16.10& 0.03& 0.03& 0.04& 1.50& 12.21& 12.21& 12.21& & 12.98& 12.84& 12.98\\ -16.30& 3.98& 3.98& 3.98& -16.00& 0.03& 0.03& 0.03& 1.52& 12.21& 12.21& 12.21& & -- & -- & -- \\ -16.22& 3.72& 3.98& 3.72& -15.90& 0.03& 0.03& 0.03& 1.54& -- & -- & -- & & -- & -- & -- \\ -16.14& -- & -- & -- & -15.80& 0.03& 0.03& 0.03& 1.56& -- & -- & -- & & -- & -- & -- \\ -16.06& -- & -- & -- & -15.70& 0.03& 0.03& 0.03& 1.57& -- & -- & -- & & -- & -- & -- \\ -15.98& -- & -- & -- & -15.60& 0.03& 0.03& 0.03& 1.59& -- & -- & -- & & -- & -- & -- \\ -15.90& -- & -- & -- & -15.50& -- & -- & -- & 1.61& -- & -- & -- & & -- & -- & -- \\ -15.82& -- & -- & -- & -15.40& -- & -- & -- & 1.63& -- & -- & -- & & -- & -- & -- \\ -15.74& -- & -- & -- & -15.30& -- & -- & -- & 1.65& -- & -- & -- & & -- & -- & -- \\ -15.66& -- & -- & -- & -15.20& -- & -- & -- & 1.66& -- & -- & -- & & -- & -- & -- \\ -15.58& -- & -- & -- & -15.10& -- & -- & -- & 1.68& -- & -- & -- & & -- & -- & -- \\ \hline \end{tabular} \end{minipage} \end{table*} \begin{table*} \label{calibtable2} \centering \begin{minipage}{170mm} \caption{Calibrating values for distance determination for the Galactic $\Sigma-D$ relations (SNRs and PNe) and the $PL$ relation (in the form of apparent magnitude $m$ vs. period $P$) for fundamental mode Cepheids in I and V band. The units of the $\log$ values and coordinate range are the same as in Figures \ref{pdf_snrs_pne} and \ref{ceph_pdf_I_V}. The grid size is $10^3 \times 10^3$. The complete table is available as an on-line material.} \begin{tabular}{ccccccccccccc@{\hspace*{-0.1mm}}ccc} \hline \multicolumn{4}{c}{--------------- SNRs ---------------}&\multicolumn{4}{c}{--------------- PNe ---------------}&\multicolumn{8}{c}{----------------------------- Cepheids -----------------------------}\\ $\log{\Sigma}$ &\multicolumn{3}{c}{$D[pc]$}&$\log{\Sigma}$&\multicolumn{3}{c}{$D[pc]$}&$\log{P}$&\multicolumn{3}{c}{$m_{\mathrm{I}}$}&&\multicolumn{3}{c}{$m_{\mathrm{V}}$}\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} &mode&mean&med.&&mode&mean&med.&&mode&mean&med.&&mode&mean&med.\\\cline{2-4}\cline{6-8}\cline{10-12}\cline{14-16} \hline -23.492& -- & -- & -- & -24.990& -- & -- & -- & -0.0982& -- & -- & -- & & -- & -- & -- \\ -23.484& -- & -- & -- & -24.980& -- & -- & -- & -0.0964& -- & -- & -- & & -- & -- & -- \\ -23.476& -- & -- & -- & -24.970& -- & -- & -- & -0.0946& -- & -- & -- & & -- & -- & -- \\ -23.468& -- & -- & -- & -24.960& -- & -- & -- & -0.0928& -- & -- & -- & & -- & -- & -- \\ -23.460& -- & -- & -- & -24.950& -- & -- & -- & -0.0910& -- & -- & -- & & -- & -- & -- \\ -23.452& -- & -- & -- & -24.940& -- & -- & -- & -0.0892& -- & -- & -- & & -- & -- & -- \\ -23.444& -- & -- & -- & -24.930& -- & -- & -- & -0.0874& -- & -- & -- & & -- & -- & -- \\ -23.436& -- & -- & -- & -24.920& -- & -- & -- & -0.0856& -- & -- & -- & & -- & -- & -- \\ -23.428& -- & -- & -- & -24.910& -- & -- & -- & -0.0838& -- & -- & -- & & -- & -- & -- \\ \hline \end{tabular} \end{minipage} \end{table*} | 14 | 3 | 1403.1368 |
1403 | 1403.4251_arXiv.txt | We compare the results from several sets of cosmological simulations of cosmic reionization, produced under Cosmic Reionization On Computers (CROC) project, with existing observational data on the high-redshift \lya\ forest and the abundance of \lya\ emitters. We find good consistency with the observational measurements and the previous simulation work. By virtue of having several independent realizations for each set of numerical parameters, we are able to explore the effect of cosmic variance on observable quantities. One unexpected conclusion we are forced into is that cosmic variance is unusually large at $z>6$, with both our simulations and, most likely, observational measurements are still not fully converged for even such basic quantities as the average Gunn-Peterson optical depth or the volume-weighted neutral fraction. We also find that reionization has little effect on the early galaxies or on global cosmic star formation history, because galaxies whose gas content is affected by photoionization contain no molecular (i.e.\ star-forming) gas in the first place. In particular, measurements of the faint end of the galaxy luminosity function by JWST are unlikely to provide a useful constraint on reionization. | \label{sec:intro} If cosmic reionization can be called the current frontier of extragalactic astronomy, then, in historic terms, we live in the middle of XIX century. I.e., the frontier is being settled... Ultra Deep Field campaigns with the Hubble Space Telescope pushed the search for the most likely reionization sources - young star-forming galaxies - to double digit values of cosmic redshift \citep{gals:biff07, gals:biol11,rei:obig12,rei:btos12,rei:sreo13,rei:wmhb13,rei:obil13,rei:bdm14,rei:obi14}. Observations of \lya\ emitters at $z\sim7$ \citep{lae:hcb10,lae:osf10,lae:pfv11,lae:ksm11,lae:sse12,lae:oom12,lae:cbw12,lae:cbw14} indicate rapid change in their abundance as one rides deeper into the frontier territory. Recent mind-blowing progress of the first generation experiments for detecting the redshifted 21cm signal from the epoch of reionization \citep{rei:pla13,rei:dlw13} promises a major observational breakthrough well before the end of this decade. Even along the well-trodden ``Oregon Trail'' of \lya\ absorption spectroscopy of high redshift quasars new advances are expected in the nearest future, as new discoveries of $z>6$ quasars continue \citep{rei:bvm13,rei:vfs14}. Theoretical studies did not stay behind the observational strides, rejuvenating a somewhat slowed-down progress of the second half of the last decade. The major push on the theory side was galvanized by the pioneering idea of \citet{reisam:fhz04}, who realized that the standard lore of large-scale structure theory, Excursion Set formalism, can be applied to studying the reionization process. That idea generated a large following of semi-analytical and semi-numerical approaches for modeling reionization \citep{reisam:fo05,reisam:fmh06,reisam:mf07,reisam:aa07,reisam:zmm11,reisam:mfc11,reisam:aa12,reisam:zgl13,reisam:btc13,ng:kg13,reisam:sm14}, and more traditional models were pursued as well \citep{reisam:cf05,reisam:cf06,reisam:sv08,reisam:mcf11,reisam:vb11,reisam:mcf12,reisam:kf12,reisam:rfs13}. Unfortunately, on the numerical simulation front the progress was less dramatic, although important advances in the simulation technology did take place \citep[e.g.][for a complete review see \protect\citet{ng:gt09}]{rei:imp06,rei:zlmd07,rei:mlz07,rei:tcl08,rei:stc08,rei:ca08,rei:lcg08,rei:ipm09,rei:at10,rei:fma11,rei:ais12,rei:sim12}. However, the primary brake on the simulation progress - insufficient computing power - is finally being released, thanks to Moore's Law. Modern High Performance Computing platforms have crossed an important threshold of ``sustained peta-scale'' performance. This level of performance, currently available on about a dozen or so (non-classified) supercomputers across the globe, offers a unique opportunity for reionization theorists to make a substantial breakthrough in our ability to model cosmic reionization with high physical fidelity, and some of the most recent simulation work already took advantage of that opportunity \citep{rei:ima14,newrei:snr14,newrei:nrs14,newrei:hdp14}. Cosmic Reionization On Computers (CROC) project is another effort in producing peta-scale simulations of reionization in sufficiently large volumes (above $100\dim{Mpc}$ in comoving units), with spatial resolution reaching down to $100\dim{pc}$, and including most (if not all) of the relevant physical processes, from star formation and feedback to radiative transfer. In the first paper in the series \citep[][hereafter Paper I]{ng:g14a} we described in complete detail the simulation design and the calibration of numerical parameters. In this paper we explore the overall process of cosmic reionization as captured by CROC simulations, and compare our theoretical predictions to several observational constraints. We deliberately limit the scope of this paper to relatively easily computable quantities, which give only a global, broad-brush view of reionization, due to the limited human effort available for the analysis of the rich, but complex simulation data. We intend to continue this paper series as more detailed, labor-intensive analysis gets completed. | \label{sec:con} We present reionization history and global characteristics of the reionization process from a suite of recent numerical simulations performed as part of the Cosmic Reionization On Computers (CROC) project. CROC simulations reproduce the observed evolution of the galaxy UV luminosity function between $z=10$ and $z=6$ well, and, hence, include realistic treatment of the dominant class of ionizing sources. We find that, in order to match the observational constraints on the post-reionization \lya\ forest at $5<z<6$, we need to set the ionizing emissivity parameter $\euv$ (that measures the escape fraction up to the resolution limit of our simulations) to just under $\euv=0.2$. However, as we also emphasized in Paper I, cosmic variance increases sharply with redshift, and at $z>6$ our simulations do not yet converge on the global properties of the IGM, such as the mean Gunn-Peterson optical depth or the volume weighted $\HI$ fraction. Since the statistical power of our simulations is much higher than the statistical reach of the existing absorption spectra of high-redshift quasars, we conclude that, unfortunately, the observations are unlikely to have reached the convergence either. In a further illustration of this, we show that the distribution of the Gunn-Peterson optical depth over the redshift intervals $\Delta z\approx 0.15$ is extraordinary wide at $z>6$, but even at $z<6$ the $\tau_{\rm GP}$ distribution retains a relatively long tail towards high values. The distributions of ionized and neutral bubbles during most of cosmic reionization is approximately flat, meaning that it is roughly equally likely for a random place of the universe to be in a large or a small bubble. We find good numerical convergence in bubble sizes down to $z\sim7$, at which point the finite sizes of our simulation boxes start biasing the distribution of ionized bubbles. That result illustrates the importance of achieving consistent numerical resolution between the gas dynamic solver and the radiative transfer solver - the mismatch between the two resolution likely results in erroneous over-propagation of ionizing radiation beyond the few mean free path lengths. We show that the equivalent width of the damping wing of \lya\ absorption increases rapidly from mellow values of $\dew\sim100\dim{km/s}$ at $z=6$ to whopping $\dew\sim2000\dim{km/s}$ by $z=7$. While $\dew$ serves only as a rough proxy for the suppression of galaxy \lya\ emission line by the neutral IGM in front of it, this result is generally consistent with the observed sharp decline in the fraction of \lya\ emitting galaxies at $z=7$ as compared to $z=6$. We also confirm conclusions from the previous simulation and analytical work that such suppression corresponds to substantial, but not dominant volume weighted neutral fraction of about 0.2. While our results on the reionization history are in good agreement with most of prior studies, we find little back reaction of reionization on the properties of early galaxies. Because galaxies that are affected by photoionization contain little molecular gas (and, hence, star formation), we find that the global star formation history is insensitive to the reionization history, i.e.\ the ``Barkana \& Loeb'' effect does not exist. A more subtle effect of reionization is in modifying the faint end slope of the galaxy UV luminosity function, but such a modification is rather small (change in the slope of about 0.1 for a unit shift in the redshift of reionization). Since predicting the faint end slope to such precision theoretically would be extremely challenging, we conclude that, unfortunately, measuring the faint end slope by JWST will not be a useful constraint on reionization, contrary to expectations \citep{misc:gmc06}. One observational constraint that we have ignored so far is the optical depth to Thompson scattering from the CMB observations by the \emph{WMAP} mission. While the history of \emph{WMAP} measurements of the Thompson optical depth is rocky, the latest value from the 9-year \emph{WMAP} data is $0.089\pm0.014$ \citep[or $0.081\pm0.012$ if other data are included in a joint fit,][]{wmap9a,wmap9b}. The value we get for fiducial sets B20.uv2 and B40.uv1 is $0.052\pm0.003$, and in the set B40.uv2 that reionizes earlier, the value for the Thompson optical depth only rises to $0.057\pm0.004$, which are only marginally (at $2\sigma$ level) consistent with the \emph{WMAP} values. A large portion, if not all, of this discrepancy is due to incomplete numerical convergence of our simulations. In Paper I we compared our fiducial runs (an equivalent of $512^3$ particles in a $20h^{-1}\dim{Mpc}$ box) with a single higher mass resolution run B20HR.uv2 that we were able to complete (an equivalent of $1024^3$ particles in a $20h^{-1}\dim{Mpc}$ box). While numerical converge tests indicate that our fiducial runs account for 55\% of all ionizing photons, the higher reslution B20HR.uv2 run accounts for 80\% of them. As the result, the Thompson optical depth raises to 0.067 in that run. Simple linear extrapolation to the limit of 100\% of ionizing radiation gives a value of 0.08 for the Thompson optical depth, fully consistent with the current observational measurements. Whether incomplete numerical convergence is, indeed, a full story will have to wait for more powerful computers, however, as at present we are unable to run the whole ensemble of higher mass resolution simulations - for example, a higher mass resolution equivalent of our planned $80h^{-1}\dim{Mpc}$ run would have $4096^3$ particles and will require of order of 200 million CPU hours, the amount not currently feasible to obtain for this kind of work. | 14 | 3 | 1403.4251 |
1403 | 1403.6915_arXiv.txt | We investigate the effects of using the full waveform (FWF) over the conventional restricted waveform (RWF) of the inspiral signal from a coalescing compact binary system in extracting the parameters of the source, using a global network of second generation interferometric detectors. We study a hypothetical population of (1.4-10)$M_\odot$ neutron star-black hole (NS-BH) binaries (uniformly distributed and oriented in the sky) by employing the full post-Newtonian waveforms, which not only include contributions from various harmonics other than the dominant one (quadrupolar mode) but also the post-Newtonian amplitude corrections associated with each harmonic, of the inspiral signal expected from this system. It is expected that the GW detector network consisting of the two LIGO detectors and a Virgo detector will be joined by KAGRA (a Japanese detector) and by proposed LIGO-India. We study the problem of parameter estimation with all 16 possible detector configurations. Comparing medians of error distributions obtained using FWFs with those obtained using RWFs (which only include contributions from the dominant harmonic with Newtonian amplitude) we find that the measurement accuracies for luminosity distance and the cosine of the inclination angle improve almost by a factor of 1.5-2 depending upon the network under consideration. We find that this improvement can be attributed to the presence of additional inclination angle dependent terms, which appear in the amplitude corrections to various harmonics, which break the strong degeneracy between the luminosity distance and inclination angle. Although the use of FWF does not improve the source localization accuracy much, the global network consisting of five detectors will improve the source localization accuracy by a factor of 4 as compared to the estimates using a three-detector LIGO-Virgo network for the same waveform model. | \label{sec:intro} Coalescing compact binary (CCB) systems, composed of NSs and/or stellar mass BHs, are among the prime targets for the second generation of GW detectors such as advanced LIGO \cite{aligo} and advanced Virgo \cite{avirgo}. On the other hand, the proposed space-based detector eLISA \cite{2012arXiv1201.3621A} shall be primarily looking at GW signals from super massive BHs. In addition, although there are no observational evidences for the existence of CCBs with intermediate mass BHs (with masses of few tens to few hundred solar masses), if at all such systems exist they should be observed by advanced ground-based detectors (see \cite{lrr-2011-5} for a review on detection of GW sources from ground and space).\footnote{We do not have observational evidence even for CCBs with stellar mass BHs but the models related to the formation of stellar mass black holes in close binaries are well supported by stellar evolution models (see for example \cite{Fryer:2001fv}).} The GW observation of stellar/intermediate mass CCB systems in advanced GW detectors will not only provide the first direct evidence for the existence of GWs but also will reveal a great deal of information about the source properties which cannot be accessed through conventional electromagnetic observations. Hence, apart from the problem of detection one is interested in estimating the parameters which characterize the source. In the case of ground-based detectors, in general one would have a situation when the GW signal is completely buried in the noise. Hence, in order to be able to detect or to extract parameters of the source one employs data analysis techniques such as matched filtering \cite{Helstrom68, Th300, Schutz89}, which in turn requires accurate modeling of the dynamics of sources emitting the signal. This has led to the development of many analytical and numerical techniques which are used to model various stages of CCB evolution, namely, the early inspiral phase, the late inspiral, the merger phase and the final ringdown phase. For instance, the early inspiral phase can be very well modeled using approximation schemes in General Relativity (GR) such as the post-Newtonian (PN) approximation~\cite{Bliving}. The late inspiral and merger phase can be computed by using Numerical Relativity~\cite{Pretorius07Review} whereas the final ringdown phase can be accurately modeled using black hole perturbation theory~\cite{TSLivRev03}. Although, in general it is believed that at the time of their formation all CCB systems possess eccentric orbits, it is reasonable to assume that in late stages of their evolution (this is precisely the stage when signals would be visible in earth-bound detectors), their orbits would become circular due to radiation reaction \cite{Pe64}. During this phase the signal from a nonspinning CCB can be approximated by a template whose frequency and amplitude steadily increases until the last stable orbit is reached. The phase~\cite{BDIWW95,BFIJ02,BDEI04} and amplitude~\cite{BIWW96,ABIQ04,KBI07,BFIS08} of GW signals from CCBs in this stage has been computed to very high accuracies using the post-Newtonian approximations in GR. Further, the fact that the amplitude of the signal in this phase varies much more slowly as compared to the phase of the signal and also because most of the signal power is contained in the dominant harmonic (quadrupolar mode), it seems reasonable to approximate the signal to a template which neglects contributions from harmonics other than the dominant one and various post-Newtonian amplitude corrections associated with each harmonic. A waveform obtained in this fashion is called the restricted waveform (RWF) and contains only the dominant harmonic at twice the orbital frequency, with phase which includes all the PN corrections to the leading phase term but only the Newtonian amplitude. Note that modes other than the dominant one are suppressed as they contribute to the waveform at a higher post-Newtonian order. In the light of this argument we refer these additional modes together with the higher order post-Newtonian corrections to the amplitude of the dominant harmonics as subdominant modes and would follow this terminology in rest of the paper. It has been argued in a number of investigations that RWFs are good enough as far as the detection of low mass binaries (M$<$10$M_\odot$) are concerned (see e.g. \cite{ChrisAnand06}). Even CCBs as massive as 25$M_\odot$ can be detected using template bank constructed using restricted waveform approximation of the inspiral signals, however, the efficiency of extracting parameters reduces as the mass of the binary increases (see the discussion in Ref.~\cite{Farr:2009pg}). It was discussed in the case of single ground-based detectors \cite{Chris06, ChrisAnand06} and in the case of space-based detector LISA \cite{AISS07} that the mass reach of GW detectors can be significantly increased by including contributions from subdominant harmonics. Such a waveform which includes contributions from various subdominant harmonics and the post-Newtonian amplitude corrections associated with each harmonic is refer to as the full waveform (FWF).\footnote{In some places in the literature the term FWF is used for inspiral-merger-ringdown waveforms. Here we simply call such waveforms as IMR waveforms or complete waveforms and reserve the term FWF for inspiral waveforms including the contributions from the subdominant harmonics and amplitude corrections associated with different harmonics.} Although, as we move towards the higher mass end, even subdominant harmonics fail to penetrate the frequency band where the detector is most sensitive. In that case it becomes important to include the contributions from the merger phase of the binary evolution. A recent work by Capano {\it et al}.~\cite{Capano:2013raa} suggests that, inspiral-merger-ringdown (IMR) waveforms based on just the contributions from dominant harmonic will be sufficient for detecting signals from binaries with total mass up to 360$M_\odot$. However, it was also mentioned that, for systems with total mass $>100M_\odot$ and with mass ratios $>4$, indeed the sensitivity of the search improves if the waveform includes contributions from subdominant modes . Another recent study based purely on numerical waveforms \cite{Pekowsky:2012sr} suggests that with the inclusion of subdominant modes of the waveform the detection volume can be significantly increased (by about 30\%) as compared to what could be achieved by using waveforms based on the RWF approximation of the inspiral signal. Some of the previous studies showed that that the inclusion of subdominant modes in the model of the GW signal not only improves the mass-reach and the detection rates of future GW detectors but also provides a more powerful template to match with the signal in order to extract the parameters of the source accurately in context of single ground based detectors \cite{SinVecc00a, ChrisAnand06b, Littenberg:2012uj} and in the case of space-based LISA \cite{SinVecc00b, MH02, HM03, AISSV07, TriasSintes07,PorterCornish08} for nonspinning binaries (see also Ref.~\cite{KKS95,AISS05} which use RWF to investigate the quality of parameter estimation). Effects of the use of FWF over RWF on parameter estimation for {\it precessing} binaries was discussed in a recent paper by O'Shaughnessy et al. \cite{O'Shaughnessy:2014dka}, where they show how the inclusion of subdominant modes improves the parameter estimation for precessing NS-BH systems observed in next generation of ground-based GW detectors. This is possible as the FWF, by the virtue of contributions from subdominant modes, has a great deal of structure, which enables one to extract parameters of the source more efficiently as compared to the case when RWF is used (see Ref.~\cite{ChrisAnand06b} for a discussion). Further, since the inclusion of subdominant modes in the waveform brings explicit dependences on the inclination angle of the binary, the degeneracy between the inclination angle and the distance of the source, which persists in the case of RWF, finally breaks. This leads to better measurement of the inclination angle of the source. Since inclination angle and distance are strongly correlated, an improvement in the measurement of the inclination angle further improves the distance measurement. In addition, as we shall see below, with FWF the polarization angle measurement also improves. This together with the inclination angle measurement enables one to constrain the orientation of the binary significantly. Since in the future we shall have a network of five ground based detectors, one can analyze the data from different detectors coherently \cite{Pai:2000zt}. Such an analysis shall not only enable one to have larger detection volume but also help one to estimate the parameters of the sources much more accurately as compared to the accuracies that can be achieved using the single detector data. Most importantly, networks with three or more detectors will be able to localize the source very accurately, which is of great importance to astrophysics and fundamental physics (see \cite{Schutz:2011tw} for a detailed discussion). The problem of parameter estimation in context of the future network of ground based detectors has been studied extensively in the past \cite{JK94, JKKT96, Ajith:2009fz, 2010PhRvD..81h2001W, Nissanke:2011ax, Klimenko:2011hz, Fairhurst:2012tf, Schutz:2011tw, Sathya.LIGOIndia}. All of these studies used RWF approximation of the GW signal to show how a network of three or more detectors shall improve the localization (or in general the measurements of parameters of the source) of the CCB system observed in the earth-bound detectors. However, Rover {\it et al}.~\cite{Rover:2006bb} considered a network consisting the initial LIGO detectors and the Virgo and investigated the accuracies with which parameters of a BNS system can be measured. They used inspiral waveforms with 2PN amplitude and phase up to 2.5PN order and used their Markov chain Monte Carlo (MCMC) routine for coherent parameter estimation. Recently, the effect of higher signal harmonics on parameter estimation of a BH-NS system was investigated in \cite{Cho:2012ed, O'Shaughnessy:2013vma} in context of a fiducial (idealized) network of two interferometric detectors using an effective Fisher matrix approach introduced in \cite{Cho:2012ed}. In this work we aim to study the effects of using the FWF over RWF on the parameter estimation for a typical nonspinning CCB system, in context future GW interferometric detectors using the Fisher information matrix approach \cite{Finn92,FinnCh93}. For this purpose we consider a population of NS-BH systems (with component masses as (1.4, 10 M$_\odot$)), all placed at a luminosity distance of 200 Mpc and distributed uniformly over the sky surface. We run simulations for about 12800 realizations obtained by randomly choosing the angular parameters giving the location and orientation of the binary. We make use of an inspiral waveform which includes amplitude corrections to various harmonics consistent up to 2.5PN order and phasing up to 3.5PN order \cite{ABIQ04}.\footnote{Inspiral waveforms with amplitude corrections to various harmonics consistent up to 3PN order are already available \cite{BFIS08} but in the present study we chose to work with a waveform which is 2.5PN accurate in amplitude.} Since it is convenient to use the waveforms in frequency domain in the Fisher information matrix approach, we use the frequency domain waveform obtained with the stationary phase approximation \cite{Th300} of the Fourier transformation of the time domain waveform of \cite{ABIQ04}. This was already computed in \cite{ChrisAnand06} and here we just use the waveform obtained there. The organization of the paper is as follows. In Sec.~\ref{sec:pe}, we first discuss the future network of advanced detectors along with the noise curves for individual detectors used in the present study. Next, we introduce our waveform model and discuss various coordinate frames which have been chosen to obtain the response of the each detector of the network. We discuss briefly our parameter estimation strategy which broadly includes the details of Fisher matrix formalism. Finally we close this section by providing the details of the system under investigation and other analysis details. In Sec.~\ref{sec:results} we list main features of the improvement in parameter estimation due to the use of FWF and compare the results for various multidetector networks. We have added a subsection to address the implications of including the LIGO-India in the global network of detectors. Finally, in Sec.~\ref{sec:discn} we summarize our results and give some future directions. | \label{sec:discn} In this paper we presented our findings of the parameter estimation study which was performed considering a population of NS-BH systems in context of the network of future advanced detectors. For the analysis we used 12800 realizations of the source (with fixed component masses of 1.4 and 10 M$_\odot$), obtained by randomizing all four angular parameters giving location ($\theta$, $\phi$) and orientation ($\iota$, $\psi$), all at a fixed luminosity distance of 200 Mpc. Our prime focus in this paper has been to investigate the quality of parameter estimation that can be achieved using amplitude-corrected waveform of inspiral signal from a nonspinning NS-BH system. For this purpose we use a post-Newtonian waveform that is 2.5PN accurate in amplitude and 3.5PN in phase given in \cite{ABIQ04}. Such a waveform is characterized in terms of nine parameters given in Eq.~\eqref{eq:param}. We use the Fisher information matrix approach to estimate all parameters of the source (see Sec.~\ref{subsec:errest} for the discussion). Our findings have been presented in Sec.~\ref{sec:results}. We discuss our results in three different subsections. In Sec.~\ref{subsec:RWF&FWF} we compare the accuracies with which various parameters of the source can be measured using both the RWF and FWF approximation to the inspiral signal mainly in context of three representative networks, namely, the LIGO-Virgo network (LHV), the LIGO-Virgo-KAGRA network (LHVK) and the LIGO-Virgo-KAGRA network after including LIGO-India (LHVKI). Although the median of the error distributions associated with each parameter for all 16 possible combination of three, four, and five detectors has been displayed in Table~\ref{tab:medianerror}, we find that for a given network the use of the FWF in general improves the parameter estimation for various parameters. However, the effect is more prominent in case of four parameters, namely, the distance ($D_L$), the inclination angle of the binary ($\cos(\iota)$), the polarization angle ($\psi$) and the phase at the coalescence epoch ($\Phi_c$). The related error distributions have been presented in Figs.~\ref{fig:compRvsFLHV}-\ref{fig:compRvFLHVKI}. Upon comparing the median errors displayed in figures as well as in Table~\ref{tab:medianerror} we find that, given the network under consideration, the errors in $D_L$ and $\cos(\iota)$ improve roughly by a factor of $1.5-2$ whereas those related to $\psi$ and $\Phi_c$ improve roughly by a factor of $1.2-1.6$. We also notice that the factor of improvement is larger for detector networks with fewer detectors. For instance, the factor of improvement in the LHV case reduces from the value of about 2 to about 1.5 for LHVKI case. This trend is in general true for all parameters. This is not very surprising as the inclusion of additional detector sites breaks the degeneracy in angular parameters such as $\psi$ and $\Phi_c$, which in turn improves the error estimation even for the RWF case, diluting the importance of the use of FWF. Measurement of other parameters does not quite improve with the use of the FWF (see Table~\ref{tab:medianerror} and the discussion presented in related subsection). In Sec.~\ref{subsec:compNetwork} we compare our parameter estimation results obtained using the FWF for three representative networks (LHV, LHVK, LHVKI). As mentioned in the beginning of Sec.~\ref{subsec:compNetwork}, although the choice of these networks for displaying our main results is mainly based on the time line argument that when various detectors would start operating, we find that they can indeed be chosen as representatives of the three, four, and five-detector networks. As should be clear from Figs.~\ref{fig:compNetwork1}-\ref{fig:compNetwork2} and the median errors displayed there, although in general the parameter estimation improves for all parameters when we add KAGRA and LIGO-India to the LIGO-Virgo network, the improvement is most significant in the case of angular resolution. The angular resolution improves almost by a factor of 2.5 with the addition of KAGRA to the LHV network where as the same improves almost by a factor of 4.5 when LIGO-India is added to the LHVK network. Again we refer to Table~\ref{tab:medianerror} for comparing the parameter estimation accuracies for all 16 possible combinations of three, four, and five-detector networks. Finally, in Sec.~\ref{subsec:LI}, we discuss in particular the benefits of adding the LIGO-India detector to the LHVK network. In addition to our conclusions based on comparisons of different networks presented in Sec.~\ref{subsec:compNetwork}, in this section we basically argue how the addition of LIGO-India detector would help achieving scientific objectives. Table~\ref{tab:medianerror-fixedsnr} corresponds to a case when the errors listed in Table~\ref{tab:medianerror} has been rescaled so that all errors would correspond to a SNR of 20. The reasons behind displaying such a table are manyfold. First and foremost it helps us quantifying various effects which play an important role in the measurement of various parameters apart from the SNR. For instance, after comparing the FWF and RWF numbers for $D_L$ and $\cos(\iota)$ errors in the two tables we find that the improvement is actually coming from the fact that the use of FWF helps breaking the $D_L$-$\iota$ degeneracy which persists in the case of RWF and SNR indeed plays no role here. Similarly, it also helps in quantify effects of having a detector network with larger areas while comparing different networks. The other reason for including the table is related to the fact that although different networks have different distance reach in our main analysis we choose to keep the sources at 200 Mpc for all network configurations. Ideally one should keep sources at different distances for different detector networks as the horizon distance for each network is different. By fixing the SNR, this issue is automatically resolved since for networks with larger horizon distance, the errors would be rescaled to values that actually correspond to the source at larger distance and vice versa. Finally, in practice sources will probably be observed with an SNR of about 20 or so. The errors displayed in Table~\ref{tab:medianerror-fixedsnr} present a more realistic scenario, which we might witness in the coming years of GW astronomy. In Table~\ref{tab:medianerror_allLIGO} we show the median errors with hypothetical detector networks in the case when all of the detector noise power spectrum is given by that of advanced LIGO and all sources are located at 200 Mpc. In Table~\ref{tab:medianerror_allLIGO-fixedsnr}, we also show the median errors with hypothetical detector networks in the case when all of the detector noise power spectrum is given by that of advanced LIGO, and the SNR is rescaled to 20. Differences between Table~\ref{tab:medianerror} and \ref{tab:medianerror_allLIGO}, and between Table~\ref{tab:medianerror-fixedsnr} and \ref{tab:medianerror_allLIGO-fixedsnr} are caused by the differences in the noise power spectrum of Virgo and KAGRA. In Sec.~\ref{subsec:compNetwork}, by comparing Table~\ref{tab:medianerror-fixedsnr} and \ref{tab:medianerror_allLIGO-fixedsnr}, we found that some unusual trends of the median errors in the mass parameters, ${\cal M}_c$ and $\delta$, in Table~\ref{tab:medianerror-fixedsnr} were caused by the differences in the noise power spectrum of LIGO, Virgo, and KAGRA. Although the Fisher analysis can be used to get a fair idea about the quality of the parameter estimation that can be achieved in future, it assumes ideal situations (such as the use of Gaussian noise) and merely provides the lower bound on errors with which various parameters can be measured. Moreover, the method is limited to the signals of high strength. In order to have a more realistic estimate of parameters of the GW source, one has to perform more realistic simulations, such as those based on Bayesian inference with real data, which are applicable to signals with arbitrary strength. However, such methods are quite expensive especially since one has to repeat the exercise for different noise realizations. Proposed variants of the Fisher matrix, such as effective Fisher matrix~\cite{Cho:2012ed}, can also be used to carry out similar studies. In addition, Ref.~\cite{Vallisneri:2011ts} provides a semianalytical technique to perform parameter estimation for signals of arbitrary strength. One can expect that this approach might be computationally bit cheaper, but an actual analysis based on this proposal is yet to be made. Besides these limitations of the Fisher analysis, the importance of the effect of abrupt termination of the waveform at $f_{\rm LSO}$ was pointed out recently in \cite{MBOFF2014}. The LSO frequency of gravitational waves is given as $1/(6^{3/2}\pi M)$ for RWF and $k/(6^{3/2}\pi M)$ for the $k$th harmonic mode of FWF, respectively. Since these values depend on the total mass, when we approximate the likelihood function by using the Taylor expansion around the true value of parameters, we have to take into account the dependence of $f_{\rm LSO}$ on the total mass. In this paper, we have not taken into account such an effect. However, as discussed in \cite{MBOFF2014}, if the detector's noise dominates the signal at the frequency of $f_{\rm LSO}$, this effect can be neglected. In Fig. 3 of \cite{MBOFF2014}, they compare the statistical uncertainty of the chirp mass $\sigma_{M_c}/M_c$ and the systematic bias, $\Delta{\hat{M}_c}/M_c$, produced by the mass-dependent LSO frequency. They consider the case of RWF, $m_1=1.35M_\odot$ and $m_2= 5-20M_\odot$, and the advanced LIGO noise spectrum which is the same as this paper. They found that, in the case of signal-to-noise ratio of 10, the systematic bias due to $f_{\rm LSO}$ dominates the statistical uncertainty if $m_2\gtrsim 11M_\odot$. In the case of the source distance of 200Mpc in our simulation, the average signal-to-noise ratio at three LIGO detectors is around $12\sim 13$ which is similar to the value. Thus, the effect of the cutoff at $f_{\rm LSO}$ might marginally affect the value of the parameter estimation errors in this paper. We expect that, as far as we are comparing the cases for FWF and RWF, and are comparing combinations of various detectors, the effect of the cutoff at $f_{\rm LSO}$ will not change the trend we observed in this paper, since the cutoff at $f_{\rm LSO}$ may affect the results of all cases in a similar way. Nevertheless, in order to obtain the definite answer to this, we need to investigate the effect of $f_{\rm LSO}$ for FWF and for the case of the network of detectors. Finally, we want to point out two important effects in the waveform modeling that we have not accounted for, which can significantly affect our estimates. First is the neglect of the spin effects in modeling the binary system. Though it may be safe to neglect the spin of the NS, the BH in the binary system may be spinning in which case our nonspinning waveforms are not adequate to describe such a system. If the BH spins are not aligned with respect to the orbital angular momentum axis of the binary, there can be precessional effects as well. One may want to revisit the problem accounting for the spin effects, say using the waveforms of ~\cite{ABFO08}, in future. The second effect we have completely ignored is the finite size effects related to the NS in the binary. Though formally the finite size effects are a 5PN in the phasing (1.5PN higher than our current 3.5PN accuracy), these effects may become significant towards the late stages of the inspiral~\cite{FH08,1PNTidal2011}. Inspiral-merger-ringdown waveform models which take into account the tidal effect in NS-BH binaries have been developed~\cite{Lackey:2013axa}, and the prospect of extracting equation of state parameters from the waveform is discussed~\cite{Lackey:2013axa}. The waveforms are calibrated to the results of the numerical relativity simulations~\cite{lrr-2011-6}. The tidal effects and the merger-ringdown phases are completely ignored in our analysis, and it may be worth revisiting the parameter estimation problem with the network of detectors by using the above-mentioned waveform models. | 14 | 3 | 1403.6915 |
1403 | 1403.6124_arXiv.txt | We study how feedback influences baryon infall onto galaxies using cosmological, zoom-in simulations of haloes with present mass $\mathrm{M}_{\mathrm{vir}}=6.9\times10^{11} \mathrm{M}_{\odot}$ to $1.7\times10^{12} \mathrm{M}_{\odot}$. Starting at $z=4$ from identical initial conditions, implementations of weak and strong stellar feedback produce bulge- and disc-dominated galaxies, respectively. Strong feedback favours disc formation: (1) because conversion of gas into stars is suppressed at early times, as required by abundance matching arguments, resulting in flat star formation histories and higher gas fractions; (2) because $50\%$ of the stars form {\it in situ} from recycled disc gas with angular momentum only weakly related to that of the $z=0$ dark halo; (3) because late-time gas accretion is typically an order of magnitude stronger and has higher specific angular momentum, with recycled gas dominating over primordial infall; (4) because 25--30\% of the total accreted gas is ejected entirely before $z\sim1$, removing primarily low angular momentum material which enriches the nearby inter-galactic medium. Most recycled gas roughly conserves its angular momentum, but material ejected for long times and to large radii can gain significant angular momentum before re-accretion. These processes lower galaxy formation efficiency in addition to promoting disc formation.\\ | \label{intro} In current cosmological models the formation of galaxies is connected to hierarchical clustering. Small structures form first, grow, and merge into larger objects. In this framework, galaxies form through the cooling of gas at the centers of dark matter haloes where it condenses into stars. To match the observed properties of galaxies and galaxy clusters, purely gravitational processes on their own cannot account for cosmological structure formation but gas dissipation processes have to be considered. Therefore it was suggested early on \citep{1978MNRAS.183..341W} that at high redshift gas has to be prevented from excessive cooling into dense regions possibly by feedback from massive stars \citep{1974MNRAS.169..229L, 1986ApJ...303...39D, 1991ApJ...380..320N}. Mechanisms like gaseous galactic outflows were proposed to remove potentially star forming low angular momentum material at early times during the formation of galaxies \citep[e.g.][]{2001MNRAS.321..471B}. Direct observational evidence of the last decades underlines potential impact of feedback events \citep[e.g.][]{1999ApJ...513..156M, 2001MNRAS.321..450T, 2003RMxAC..17...47H, 2006ApJ...637..648S, 2009ApJS..181..272G, 2010ApJ...717..289S, 2011ApJ...733L..16S}. Numerical simulations of galaxy formation are an important tool for understanding the impact of feedback. In simulations, the stellar components of galaxies are usually made from cooling halo gas and from gas and stars added by mergers. During gas poor stellar mergers the orbitals of stars are scrambled, the stellar systems become dynamically hot, and a significant part of the stellar angular momentum is transported outwards to the dark matter component \citep[e.g.][]{1992ApJ...393..484B, 1992ARA&A..30..705B}. To produce dynamically cold and thin stellar discs, the accretion of high angular momentum gas from outer regions of the haloes is needed in the more recent past \citep{1979Natur.281..200F}. This calls for gas reservoirs at low redshifts as well as for feedback processes at higher redshifts to avoid early over-cooling and overly efficient star formation. Some progress has been made in simulating galaxies by increasing resolution \citep{2003ApJ...597...21A, 2004ApJ...607..688G} or applying more elaborate models for the interstellar medium and star formation \citep[e.g.][]{2008ApJ...680.1083R, 2010Natur.463..203G, 2012MNRAS.425.3058C, 2012MNRAS.421.3488H, 2013MNRAS.432.2647H} or for stellar feedback \citep[e.g.][]{1999ApJ...519..501S, 2001ApJ...555L..17T, 2005MNRAS.363.1299O, 2008MNRAS.389.1137S, 2011MNRAS.410.1391A, 2012MNRAS.426..140D, 2013MNRAS.428..129S, 2013ApJ...777L..38H, 2013MNRAS.434.3142A, 2013MNRAS.436.3031V}. Commonly known problems in disc galaxy formation such as undersized disc galaxies and the significant loss of angular momentum from gas particles to the dark matter component, known as the angular momentum catastrophe \cite[e.g.][]{1991ApJ...380..320N, 1995MNRAS.275...56N, 1998MNRAS.300..773W, 2000ApJ...538..477N, 2002MNRAS.335..487M}, have been worked on for several years. By now, many of the past problems have been solved and more veritable disc galaxies can be formed \citep[e.g.][]{2009MNRAS.396..696S, 2011ApJ...742...76G, 2011MNRAS.410.2625P, 2011MNRAS.410.1391A, 2012MNRAS.424.1275B, 2013MNRAS.434.3142A, 2014MNRAS.437.1750M}. In particular the study of the evolution of the gas component in simulations will help to better understand the physical processes regulating galaxy formation. \cite{2008MNRAS.387..577O} and \cite{2010MNRAS.406.2325O} introduced the concept of `wind recycling', describing fluid elements which are ejected in a wind, then re-accreted and ejected again or, alternatively, condensed completely into stars. They tracked the gas particles during hydrodynamical SPH simulations and monitored if and how often they entered a wind mode. Most `wind' particles do not stay in the intergalactic medium (IGM) but are re-accreted onto the galactic halo possibly several times. They may never actually reach the IGM but circle within a so called `halo fountain', the name chosen as reference to the `galactic fountain', first introduced by \cite{1976ApJ...205..762S}, which describes a similar process on smaller galactic scales. The halo fountain is the dominating recycling process at late times in the simulations of Oppenheimer et al. ($z \leq 1$), when the wind particles typically remain within the parent halo. \cite{2011MNRAS.415.1051B} found in their simulations with supernova-powered outflows that during the assembly of the galaxy, low angular momentum material tends to be ejected at early times when the potential wells of the forming galaxies are still shallow. In addition, gas is primarily ejected perpendicular to the disc whereas inflow occurs mainly in the disc plane \citep[but see][]{1309.5951}. In a subsequent paper \cite{2012MNRAS.419..771B} found that the ejected gas is sometimes re-accreted at later times with additional angular momentum gained through mixing with hot corona gas. While the ejection of gas was found to be an important process at all galaxy mass scales, the redistribution of angular momentum via the re-accretion of gas in galactic fountains gets more relevant for higher mass galaxies. \citeauthor{2012MNRAS.419..771B} therefore concluded that galactic fountains may lead to the formation of high-mass disc galaxies. For a more detailed understanding of the effects of different feedback implementations a direct comparison is required, i.e. different feedback models applied to the same initial conditions. This was studied, for instance, in the Aquila comparison project by \cite{2012MNRAS.423.1726S}, who compared the outcome of simulations with identical initial conditions but 13 models with differing hydrodynamics and feedback schemes. The differences they found underline the importance of direct code comparison and might encourage further work in this direction \citep[see also e.g.][]{2005MNRAS.363.1299O, 2010MNRAS.409.1541S, 2011MNRAS.410.2625P}. In this paper, we present a study of the assembly history of baryonic matter in two sets of SPH simulations with different feedback implementations but identical initial conditions. We follow the impact of the feedback from massive stars right from the beginning of our simulations until the present day. Thereby we can study its effect on the assembly of the baryons which starts long before they are accreted onto the galaxy and in many cases may even prevent this. The paper is organized as follows: In Section~\ref{simulations} we describe the simulations and the different feedback models used for the comparison. Section~\ref{global} contains an analysis of general differences in the simulations regarding baryon conversion, angular momentum distribution and mass accretion. In Section~\ref{assembly} we present a more detailed analysis of the accretion history of gas and stars onto forming galactic discs, including a study of the angular momentum distribution of the gas. We summarize and discuss our results in Section~\ref{summary}. | \label{summary} We present a comparative study of the gas assembly histories for five cosmological galaxy zoom-in simulations ($\mathrm{M}_{\mathrm{vir}}=6.9\times10^{11} \mathrm{M}_{\odot}$ to $1.7\times10^{12} \mathrm{M}_{\odot}$) each carried out twice, once with weak (WFB), and once with strong feedback (SFB) from massive stars. We summarize characteristic quantities of those simulations in Table~\ref{tab:4}. The relative amount of stellar and gas accretion at different redshifts and the angular momentum of the accreted gas determine whether stars form and assemble in a spheroid (low angular momentum) or a disc (high angular momentum). We focus on the detailed evolutionary histories of the baryonic components since $z=4$ and track accretion onto the halos and accretion and ejection from the galaxies, as well as the full evolution of the disc gas angular momentum during ejection and accretion. In the WFB simulations (see e.g. \citealp{2010ApJ...725.2312O,2013arXiv1311.0284N}) the early ($z>1$) conversion of low angular momentum gas into stars in the galaxies and in the accreted substructures is favoured. This leads to the formation of systems with a high stellar-to-halo mass-ratio \citep{2010ApJ...725.2312O}, relatively low angular momentum, and a significant spheroidal component \citep{2012ApJ...754..115J,2013arXiv1311.0284N}. There is little late accretion of gas, which could form a (more extended) disc-like component \citep[see also][]{2014arXiv1401.3180S}. Overall, the behaviour reflects the well documented angular momentum `problem' or `catastrophe' (e.g. \citealp{2000ApJ...538..477N}). Simulations starting from identical initial conditions using a strong feedback implementation show different behaviour and the galaxies better resemble the observed population of present-day spiral galaxies \citep{2013MNRAS.434.3142A}. The early conversion of low angular momentum gas into stars is -- in agreement with abundance matching estimates -- significantly suppressed. This leads to higher gas fractions at all redshifts and to flat star formation histories with less star formation at high redshift and more star formation at low redshift. This is a common feature of models with strong stellar feedback either directly coupled to the surrounding ISM (see e.g \citealp{2010MNRAS.408..812S,2011MNRAS.410.2625P,2013MNRAS.434.3142A,2013MNRAS.428..129S}) or decoupled from the local ISM in a wind mode (see e.g \citealp{2008MNRAS.387..577O,2010MNRAS.406.2325O,2014MNRAS.437.1750M,2013MNRAS.436.2929H,2013MNRAS.436.3031V,2014MNRAS.tmp...38T}). Our quantitative results regarding feedback efficiency and galaxy morphology agree with the qualitative conclusions for high-$z$ galaxies presented in \cite{2014ApJ...782...84A}. Present-day gas is arranged in an extended disc. Most galactic stars in the SFB model form within the galaxy from disc gas, so the stellar component forms a disc and its angular momentum is a poor reflection of that of the dark matter halo. The gas accretion rates onto the central regions of the halo and the galactic disc are in general significantly higher in the SFB model (5-15 times). At early times ($z>1$) the accretion rates are dominated by first accreted gas, while the accretion of recycled gas becomes dominant at late times ($z<1$). The recycling of disc gas is an important feature of the SFB model with 50-60 per cent of the accreted gas mass participating in this process and thus forming a galactic fountain \citep{2008MNRAS.387..577O}. Also, the high gas accretion rates at lower redshift decouple from the declining dark matter accretion \cite[see][]{2013MNRAS.436.2929H}. In general, the angular momentum of first accreted as well as recycled gas increases significantly towards low redshift \cite[see also][]{2012MNRAS.419..771B}, resembling (for first accreted gas) an inside-out growth of the disc. This process might be influenced or delayed through mergers at early times. A significant fraction (25-30 per cent) of the accreted disc gas is ejected and does not return by the present day. Confirming previous studies \citep{2011MNRAS.415.1051B} this gas is predominantly ejected at high redshift ($z>1$) - when the potential wells of the proto-galaxies are still shallow - and has low angular momentum. The predominant ejection of low angular momentum gas has long been proposed as a possible solution to the angular momentum problem \citep{2001MNRAS.321..471B,2002MNRAS.335..487M,2006MNRAS.372.1525D,2009MNRAS.396..141D} and the promoting effect for the formation of discs has been confirmed by cosmological simulations \citep{2010Natur.463..203G,2011MNRAS.415.1051B}. The ejected low angular momentum gas can by $z=0$ reach distances from the galaxy of the order of $\gtrsim$~1~Mpc and can thus contribute to metal enrichment of the IGM. The efficient ejection of low angular momentum gas also prevents conversion into stars even in major galaxy mergers. Mergers can trigger angular momentum loss of gas and have long been considered a major problem for disc galaxy formation \citep[e.g.][]{1996ApJ...471..115B}. Only recently has there been growing evidence that - provided efficient star formation is suppressed and/or the feedback is sufficiently strong - gas-rich early mergers are less problematic \citep{2005ApJ...622L...9S,2006ApJ...645..986R,2009ApJ...691.1168H,2009MNRAS.398..312G,2011MNRAS.415.3750M} and are to some degree even required to explain the structural evolution of the disc galaxy population (\citealp{2014arXiv1404.6926A}; see e.g.\, \citealp{2005A&A...430..115H} for observational evidence of disc reformation in the last 8 Gyrs after major mergers). The angular momentum of cycling gas particles is in general conserved, due to short typical travel distances (less than 10 kpc for 60-85 per cent of the gas) and short recycling times ($<$~1~Gyr). Some smaller fraction ($\sim$10 per cent) of cycling gas - mostly ejected during early turbulent phases of merging - gains more than 1000 kpc km s$^{-1}$ before its re-accretion, eventually through mixing with the hot corona gas \citep{2012MNRAS.419..771B} or from cosmic torques. However, if gas particles leave the disc for long times ($>$~1~Gyr) and travel large distances ($\gtrsim$ 50-100 kpc), then they always gain angular momentum before re-accretion. However, this gas contributes little ($<$ 3 per cent) to the total gas accretion. Nonetheless, the general trend towards larger travel distances and longer recycling times in SFB as compared to WFB is strong. All the above processes resulting from strong stellar feedback favour the formation of extended galactic discs. | 14 | 3 | 1403.6124 |
1403 | 1403.2707_arXiv.txt | { Using several lines of evidence we show that the scale values of the geomagnetic variometers operating in Helsinki in the 19th century were not constant throughout the years of operation 1844-1897. Specifically, the adopted scale value of the Horizontal Force variometer appears to be too low by $\sim$30\% during the years 1866-1874.5 and the adopted scale value of the Declination variometer appears to be too low by a factor of $\sim$2 during the interval 1885.8-1887.5. Reconstructing the Heliospheric Magnetic Field strength from geomagnetic data has reached a stage where a reliable reconstruction is possible using even just a single geomagnetic data set of hourly or daily values. Before such reconstructions can be accepted as reliable, the underlying data must be calibrated correctly. It is thus mandatory that the Helsinki data be corrected. Such correction has been satisfactorily carried out and the HMF strength is now well constrained back to 1845. | 14 | 3 | 1403.2707 |
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1403 | 1403.0644_arXiv.txt | The Telescope Array is the largest experiment studying ultra-high energy cosmic rays in the northern hemisphere. The detection area of the experiment consists of an array of 507 surface detectors, and a fluorescence detector divided into three sites at the periphery. The viewing directions of the 38 fluorescence telescopes point over the air space above the surface array. In this paper, we describe a technique that we have developed for simulating the response of the array of surface detectors of the Telescope Array experiment. The two primary components of this method are (a) the generation of a detailed CORSIKA Monte Carlo simulation with all known characteristics of the data, and (b) the validation of the simulation by a direct comparison with the Telescope Array surface detector data. This technique allows us to make a very accurate calculation of the acceptance of the array. We also describe a study of systematic uncertainties in this acceptance calculation. | The Telescope Array (TA) is the largest experiment studying ultrahigh energy cosmic rays in the northern hemisphere. It is located in Millard County, Utah, and consists of a surface detector (SD) of 507 scintillation counters, each of area 3m$^2$, deployed in a grid of 1.2 km spacing, plus a set of 38 fluorescence telescopes located at three sites around the SD looking inward over the array. Both detector systems of TA started collecting data in 2008. Measurements of the differential flux of cosmic rays, as a function of energy, have historically played an important role in the study of ultra-high energy cosmic rays (UHECRs). Foremost among these is the high energy break in the spectrum at $5 \times 10^{19}$ eV, called the GZK cutoff~\cite{greisen}\cite{zk}\cite{hiresmono:gzk}\cite{auger_gzk}, which provides convincing evidence for the extra-galactic origin of the highest energy cosmic rays. An important experimental technique used in the spectrum measurement is the calculation, using the Monte Carlo simulation method, of the efficiency with which the detector observes cosmic ray induced extensive air showers. Prior to the TA experiment, high-fidelity Monte Carlo (MC) simulations have been available for fluorescence detectors (FDs), which measure the fluorescence light emitted by nitrogen molecules excited by the passage of shower particles in their vicinity. Accurate simulations for the other major detector type, surface scintillation arrays, have only recently become possible with the rapid growth of computational and storage capacity over the past decade, coupled with the maturity of sophisticated and realistic shower generation codes over the same time frame. In particular, the difficulty of generating accurate Monte Carlo simulations of air showers has limited the surface array technique to the energy regime where the detector is 100\% efficient \cite{auger:optdist}; i.e., only at the high energy end of the detector's sensitive range. In order to simulate accurately the ground-level particle densities measured by surface detectors, along with their fluctuations, a shower generator code needs in principle to track every particle created in the avalanche process down to below its critical energy. In practice, available CPU power and storage space limit one to generating only a small number of shower particles, insufficient for an accurate calculation of detector acceptance, or for a useful comparison of data and MC distributions. An approximation technique called "thinning" \cite{Hillas:1997tf} typically is used in programs like CORSIKA~\cite{corsika} and AIRES~\cite{aires} to reduce CPU time requirements. Under the thinning approximation, nearly all particles with energies below a preselected threshold (orders of magnitude higher than the critical energy) are removed from the shower. Only a few representative particles are kept with weights to account for those, in the same region of phase space, that have been "thinned" out. The thinning method usually gives an adequate description of particle distributions in the core region of a shower where enormous numbers of particles are found (and where essentially all of the fluorescence light is generated). For surface detectors, which sample the particle density at ground level, the enormous flux saturates any counter in proximity to the shower core. Typically, useful sampling is based on detectors at the scale of the detector spacing or more. For experiments, like TA, that are optimized to measure the highest energy cosmic rays, this distance scale is of the order of a kilometer. While a thinned shower is able to reproduce the average particle densities reasonably well on the kilometer scale from the shower core, the weighted particles cannot model the shower-to-shower fluctuations or even the fluctuations at different azimuthal angles around the shower core. The RMS deviations from the average densities in a thinned shower are typically off by an order of magnitude or more from that obtained from those seen in the few "unthinned" showers one can afford to generate. Thinning is therefore too crude of an approximation to give a faithful representation of even the simulated air shower itself, let alone real cosmic-ray induced showers. Some experiments have claimed to overcome this intrinsic difficulty by restricting their analysis to the highest energy range where the efficiency of the detector approaches unity. However, if quality cuts are used to select only a subset of the data, then the use of a simulation is still needed to calculate acceptance. In that case the use of thinning can and probably does introduce significant systematic biases because the thinned Monte Carlo (MC) simulation cannot accurately reproduce the tails expected in the distribution of cut parameters. Quality cuts are invariably used to remove outliers in such tails. In the simulation of air showers for calculating the acceptance of the Telescope Array experiment, we have developed a "de-thinning" procedure to compensate for the shortcomings of the thinning. Using the thinned CORSIKA output, we replace each representative particle of weight $w$ with an ensemble of $w$ particles propagated in a cone about the weighted particle. A detailed prescription of our de-thinning process was published in an earlier article~\cite{dethinned}. In that article, careful comparisons were made between de-thinned and non-thinned showers (the latter referring to showers generated without any thinning), and excellent agreement was found in the statistical properties of the two sets of simulations. Our de-thinned sample overcame all of the essential shortcomings of the thinning approximation. In this paper, we describe the actual application of the de-thinning process to the simulation of the Telescope Array experiment. Detailed comparisons are shown for key distributions (those that directly affect the acceptance calculation) between TA data and the de-thinned MC shower sample. The excellent agreement in these comparisons serve to demonstrate the high degree of accuracy of the simulation in reproducing the properties both of the detector and of the data. This paper is the last in a series of three describing simulation techniques used for the surface detector of the Telescope Array experiment~\cite{dethinned}\cite{parallel}. Sections~2 and 3 give an overview of the TA surface detector and its data analysis. Section~4 describes the process of generating de-thinned CORSIKA showers for the TA SD, with events generated according to previous measurements of the UHECR spectrum and composition, and including a detailed simulation of each scintillation counter. Validating the Monte Carlo simulation is described in Section~5, and the experimental resolution is presented in Section~6. Determining the energy scale using events seen by both the fluorescence and surface detectors is given in Section~7, and a study of systematic uncertainties and biases is described in Section~8. \section {TA Surface Detector Data} The TA surface detector has been described previously~\cite{ta:1}\cite{ta:2}\cite{sdspec}. In Figure~\ref{fig:layout}, \begin{figure}[t,b] \begin{center} \includegraphics[width=0.5\textwidth]{fig01} \end{center} \caption{The physical layout of the Telescope Array. The surface detectors are represented by small open squares. Additionally the positions and fields of view are shown for all three air-fluorescence stations.} \label{fig:layout} \end{figure} we see the physical layout of all components of TA. Each SD counter consists of two layers of plastic scintillator, 3 m$^2$ in area, and read out independently by two photomultiplier tubes. Scintillation light is guided to the photomultiplier tubes by a system of wavelength-shifting fibers set in grooves in the scintillator. These counters are calibrated every 10 minutes~\cite{tasd:performance} using a histogram of pulse heights recorded for events triggering both layers of scintillators in time coincidence. The resulting distributions typically consist of a peak at low integrated pulse area, accompanied by a tail toward higher pulse area. The peak itself corresponds to the signal from single muons passing through both scintillators, and the centroid of the peak then defines the average signal for a minimum-ionizing particle (MIP) for that channel. The waveforms from the two scintillators are sampled by a 50 MHz FADC system [14]. A real time integration process is used to trigger each counter: waveforms of pulses with integrated areas corresponding to at least 1/3 MIPs are saved with a corresponding GPS time stamp. The detection of a pulse of 3 MIPs or larger is reported to a central data acquisition system via the radio communication system. A trigger for the TA SD occurs when three adjacent counters have energy deposits equivalent to at least three MIPs in each counter within an 8 $\mu$sec window. When the SD trigger conditions are met, all counters in the array are polled. Saved waveforms (of minimum 1/3 MIPs) with time stamps within 64 microseconds of the event trigger are read out over the radio link. The pulse area contained in recorded waveforms, stored in raw FADC units, are converted to units of vertical-equivalent muons (VEM). A vertical minimum-ionizing muon deposits on average 2.05 MeV of ionization energy in each scintillator layer. The conversion process uses the calibration histograms collected every 10 minutes, but also incorporates the simulated detector response to single muons and other secondary particles produced in air showers induced by TeV cosmic rays. The application of the calibration and the signal analysis extract the following information from each counter participating in the event: (a) An integrated particle count in units of VEM, (b) the arrival time of the shower, and (c) the spatial coordinates of the counter. These quantities are then used to reconstruct the shower trajectory and the energy of the primary cosmic ray. Figure~\ref{figure:event_display} \begin{figure}[t,b] \begin{center} \includegraphics[width=0.5\textwidth]{fig02} \end{center} \caption { A typical high energy event seen by the TA SD. Each circle represents a counter that participated in the event. The area of each circle is proportional to the logarithm of the VEM signal size for that counter. The measured arrival time of the shower at each counter is denoted by the color of the circle. The arrow represents the projection of the shower axis onto the ground, $\hat{u}$, and the intersection between this arrow and its perpendicular bisector marks the location of the shower core. } \label{figure:event_display} \end{figure} shows a footprint of a typical high energy event. \section {Event Reconstruction and Selection Cuts} The event reconstruction procedures used for TA SD data are based on parametrizations and procedures originally developed by the AGASA Collaboration~\cite{Takeda:2002at}, modified to match the characteristics of the TA detectors~\cite{tasd_ichep2010}. First, the shower axis is determined from the arrival times in the triggered counters. These are fit to the AGASA-modified Linsley time delay function~\cite{Linsley:Td}\cite{agasa:timefit}. Figure~\ref{figure:timefit}, \begin{figure*}[t,b] \centering \subfloat[]{\includegraphics[width=0.5\textwidth]{fig03a}\label{figure:timefit}} \subfloat[]{\includegraphics[width=0.5\textwidth]{fig03b}\label{figure:ldffit}} \caption{ Time and lateral distribution fits for a typical TA SD event. \protect\subref{figure:timefit} Counter time versus distance from the shower core along the $\hat{u}$-direction. Points with error bars are the measured counter times. The solid curve gives the times predicted by the modified AGASA fit for the counters lying on the $\hat{u}$-axis. The dashed and dotted lines are the fit expectation times for the counters that are located 1.5 and 2.0 km off the $\hat{u}$-axis, respectively. \protect\subref{figure:ldffit} Measured lateral distribution fit to the AGASA LDF function. The vertical axis is the signal density and the horizontal axis is the lateral/perpendicular distance to the shower core. Event $S800$ is determined from the fit curve. } \label{figure:typical_event} \end{figure*} shows this fit for a typical TA SD event. For this work, the parameters of the original AGASA time delay function were adjusted to fit the overall characteristics of the TA SD data set, by considering the distributions of fit residual in distance-to-core, VEM signal sizes, and in zenith angle. It should be emphasized that the adjustments were made based exclusively on actual TA SD data without any additional information from simulations, and are therefore model-independent. The primary energy estimation of TA SD events is established by first measuring the charge density at 800m in lateral (perpendicular) distance from the shower axis ($S800$)~\cite{auger:optdist}. The measured particle densities from the counters are fit to the modified AGASA lateral distribution function (LDF)~\cite{agasa:results}, as shown in Figure~\ref{figure:ldffit}. The value at 800m, denoted as $S800$, is interpolated from this fit. In order to achieve reasonable detector resolutions in energy and pointing direction, but without losing an unreasonable fraction of events, we chose the following event selection cuts (both pre- and post- reconstruction). These same cuts are applied to both data and Monte-Carlo in the present TA SD analysis: \begin{enumerate} \item $N_{\mathrm{SD}} \ge 5$. At least 5 good counters per event. \item $\theta < 45^{\circ}$. Zenith angle less than 45 degrees. \item $D_{\mathrm{Border}} \ge 1200$~m. Core position is within the array and at least 1200~m away from the edge of the array. \item $\chi_{\mathrm{G}}^{2}/\mathrm{d.o.f.} < 4$ and $\chi_{\mathrm{LDF}}^{2}/\mathrm{d.o.f.} < 4$. Reduced values of $\chi^{2}$ of geometry ($\chi_{\mathrm{G}}^{2}/\mathrm{d.o.f.}$) and LDF ($\chi_{\mathrm{LDF}}^{2}/\mathrm{d.o.f.}$) fits are less than 4. \item $\sqrt{\sigma_{\theta}^{2} + \mathrm{sin}^{2}\theta \, \sigma_{\phi}^{2}} < 5^{\circ}$. Pointing direction uncertainty is less than 5 degrees. $\sigma_{\theta}$ and $\sigma_{\phi}$ are the uncertainties on zenith and azimuthal angles from the geometry fit. \item $\sigma_{S800} / S800 < 0.25$. Fractional uncertainty of $S800$ determination (from the LDF fit) is within 25\%. \end{enumerate} Table~\ref{table:tasd_cut_eff} displays the efficiency (fraction of events retained) when the quality cuts are applied, incrementally for 3 energy slices. \begin{table}[!htbp] \begin{center} \begin{tabular}{|l|l|l|l|l|} \hline Quality cut & Efficiency, & Efficiency, & Efficiency, \\ & $E > 10^{18}$~eV & $E > 10^{18.5}$~eV & $E > 10^{19}$~eV \\ \hline $N_{\mathrm{SD}} \ge 5$ & 0.674 & 0.931 & 0.973 \\ \hline $\theta < 45^{\circ}$ & 0.741 & 0.702 & 0.677 \\ \hline $D_{\mathrm{Border}} \ge 1200$~m & 0.865 & 0.814 & 0.748 \\ \hline $\chi_{\mathrm{G}}^{2}/\mathrm{d.o.f.} < 4, \, \chi_{\mathrm{LDF}}^{2}/\mathrm{d.o.f.} < 4$ & 0.928 & 0.938 & 0.981 \\ \hline $(\sigma_{\theta}^{2} + \mathrm{sin}^{2}\theta \, \sigma_{\phi}^{2})^{1/2} < 5^{\circ}$ & 0.656 & 0.925 & 0.995 \\ \hline $\sigma_{S800} / S800 < 0.25$ & 0.534 & 0.887 & 0.995 \\ \hline All cuts combined & 0.14 & 0.41 & 0.48 \\ \hline \end{tabular} \end{center} \caption{Efficiency of the quality cuts} \label{table:tasd_cut_eff} \end{table} \section {Surface Detector Monte Carlo Simulation} In simulating an air shower, the TA surface detector Monte Carlo uses the CORSIKA 6.960~\cite{corsika} simulation package. For the standard simulated event set, we selected the QGSJET-II-03~\cite{qgsjet} and FLUKA2008.3c~\cite{fluka1}\cite{fluka2} hadronic models for high and low energies, respectively. For electromagnetic processes, the EGS4~\cite{egs4} electromagnetic model was used. The first step in generating a comprehensive simulation of the TA SD data set is to create a library of thinned CORSIKA showers. This library consists of 16,800 extensive air showers with primary energies distributed in $\Delta \log_{10}E=0.1$ bins between $10^{16.75}$~eV and $10^{20.55}$~eV. The number of showers in each bin ranges from 1000 in the lowest energy bin to 250 in the highest energy bin. These showers are simulated with zenith angles from $0^\circ$ to $60^\circ$ assuming an isotropic distribution. It is important to note that in our final analysis we only include events with $E>10^{18.0}$~eV and $\theta<45^\circ$. However, events must be simulated well beyond these limits in energy and inclination in order to give a complete understanding of our detector acceptance as well as our energy and angular resolutions. Each shower in the CORSIKA library is then subjected to dethinning~\cite{dethinned}. For each simulated event, all shower particles that strike the ground are divided spatially by their landing spots into $6\times6{\rm m^2}$ ``tiles'' on the desert floor and into $20{\rm ns}$ wide bins by their arrival time. The total energy deposited by all particles that landed in a particular tile, and into a virtual TA SD counter located at its center, is calculated using the GEANT4 simulation package~\cite{geant4}. Note this analysis assumes many more virtual SD counters (spaced every 6 m instead of 1.2 km) than are actually present in the experiment. Back scattering of particles striking the ground within the tile is included in the simulation. The energy deposited as a function of time is stored in the shower library. Figure~\ref{fig:dethinning} \begin{figure*}[t,b] \centering \subfloat[]{\includegraphics[width=0.5\textwidth]{fig04a}\label{fig:th_v_nth}} \subfloat[]{\includegraphics[width=0.5\textwidth]{fig04b}\label{fig:dth_v_nth}} \caption{A comparison of energy deposition per counter versus perpendicular distance-to-core for a non-thinned and a thinned simulation before (left) and after (right) the dethinning procedure is applied. Both simulations are of a proton with a primary energy of $10^{19}$~eV and a primary zenith angle of $45^\circ$. While the mean energy deposition agrees in all cases, the variation in the energy deposition (RMS) shows much better agreement after dethinning.} \label{fig:dethinning} \end{figure*} shows the comparison of energy deposition in SD counters vs. distance-to-core from a simulated $10^{19}$ eV shower before and after de-thinning. The plot on the right, made using a de-thinned shower, shows excellent agreement to an identical unthinned shower in both the mean energy deposit and its RMS variation, plotted as functions of distance-to-core. In contrast, the same plot on the left comparing the same shower after thinning to the same unthinned shower shows a discrepancy in the RMS variation in energy deposition by up to an order of magnitude. In the concluding step of the shower library generation, each tiled shower is sampled ~2000 times through a detailed simulation of the detector, including electronics. The shower core positions, the azimuth of the shower axis, and event times are varied in this process. The detector simulation utilizes real-time calibration information from the TA SD to effect a highly detailed, time-specific simulation of the detector operating conditions. Additionally, random background particles are inserted into the electronics readout based on secondary flux derived from additional CORSIKA simulations of the low-energy cosmic ray spectrum reported by the BESS Collaboration~\cite{bess}. The net result of this step is to convert each dethinned CORSIKA shower into an event library of simulated detector events in a data format identical to that produced by the TA SD instrumentation. In order to achieve a highly accurate representation of the actual TA SD data set, we sample simulated events from our event library with a primary energy distribution and composition according to published HiRes energy spectrum~\cite{hiresmono:gzk} and composition~\cite{hires_comp}, respectively. The resulting MC event set is then processed by the same analysis program as the TA SD data. This process chain is illustrated by the diagram in Figure~\ref{figure:flowchart}. \begin{figure*}[t,b] \begin{center} \includegraphics[width=\textwidth]{fig05} \end{center} \caption{Steps for simulating the TA SD data set. Each box represents one or more computational routines used to produce the input files required for the next step.} \label{figure:flowchart} \end{figure*} Finally, we validate the accuracy of the simulation by comparing distributions of key observables obtained from the MC events with those from real data. As will be seen in the next section, our de-thinned shower samples give excellent agreement in these data-MC comparisons, thereby verifying the reliability of detector resolutions and of the detector acceptance calculation obtained from our simulation algorithm. The primary advantage of this process lies in that the trigger efficiency, reconstruction quality cuts, and the effects of finite resolution~\cite{unfolding}\cite{hiresmono:dtmc} are all automatically included in the analysis. A functional relationship between $S800$, the primary zenith angle ($\theta$), and the primary energy is constructed using the de-thinned Monte Carlo Event set. Each simulated event is subjected to the same geometrical reconstruction as described above, and the value of $S800$ obtained in the same way. A three-dimensional scatter plot is then made of the input (generated) primary energy of each shower plotted in the $z$-direction, vs. $\sec\theta$ in the $x$-direction, and the logarithm of the $S800$ value in the $y$-direction. The points in this plot form a surface that represents the shower energy as a bi-variate function of $\sec\theta$ and $\log_{10}(S800)$. The function obtained for this work is shown in Figure~\ref{figure:entable}, \begin{figure}[t,b] \centering \includegraphics[width=0.5\textwidth]{fig06} \caption { Energy as a function of reconstructed $S800$ and $\sec\theta$ made from the dethinned MC event set. The true (input) values of the primary energy are represented (along the z-axis) by color according to the key shown to the right.} \label{figure:entable} \end{figure} in which the value of energy is represented by color according to the key attached to the right of the plot. The information contained in Figure~\ref{figure:entable} is used to determine the energy of both real and simulated events from the interpolated $S800$ values. | We have demonstrated that the dethinned CORSIKA/QGSJET-II-03 proton Monte Carlo simulation accurately models the response of the TA array of scintillation counters to cosmic rays in the $E>10^{18.2}$eV, $\theta < 45^{o}$ domain. In reconstruction of events, fits to counter times and pulse heights are almost identical for the Monte Carlo and the data. Basic histograms of geometrical quantities, and of quantities related to the lateral distribution of counter pulse heights, for the Monte Carlo agree very well with the same distributions for the data. We have measured a 27\% correction to the energy scale of the CORSIKA + QGSJET-II simulation package based on air showers observed calorimetrically by the Telescope Array fluorescence detector, and examined some sources of systematic errors in our aperture calculation. We conclude that this Monte Carlo simulation is an accurate tool for calculating the surface detector aperture used to calculate the energy spectrum, as well as to estimate the exposure on the sky for cosmic ray anisotropy analyses. | 14 | 3 | 1403.0644 |
1403 | 1403.5850_arXiv.txt | We present deep observations of the Galactic supernova remnant IC\,443 with the {\it Suzaku X-ray satellite}. We find prominent K-shell lines from iron and nickel, together with a triangle residual at 8--10~keV, which corresponds to the energy of the radiative recombination continuum (RRC) of He-like iron. In addition, the wavy residuals have been seen at $\sim$5.1 and $\sim$5.5~keV. We confirm that the residuals show the first enhanced RRCs of He- and H-like calcium found in supernova remnants. These facts provide robust evidence for the recombining plasma. We reproduce the plasma in the 3.7--10~keV band using a recombining plasma model at the electron temperature 0.65~keV. The recombination parameter $n_{\rm e}t$ ($n_{\rm e}$ is electron density and $t$ is elapsed time after formation of a recombining plasma) and abundances of iron and nickel are strongly correlated, and hence the errors are large. On the other hand, the ratio of nickel to iron relative to the solar abundances is well constrained to 11$^{+4}_{-3}$ (1$\sigma$). A possibility is that the large abundance ratio is a result of an asymmetric explosion of the progenitor star. | Thermal X-rays from the supernova remnants (SNRs) are due to shock-heated plasma. Electrons are first heated to high temperatures ($T_{\rm e}$), and then the atoms are gradually ionized. Therefore, the plasmas in young and intermediate-aged SNRs are often described by a non-equilibrium ionization plasma (NEI), or more specifically an ionizing plasma (IP). Electron temperature ($T_{\rm e}$) is higher than that predicted from the mean ionization states of elements (ionization temperature: $T_{\rm z}$). As SNRs age, the ionization rate decreases and is balanced by the recombination rate; therefore, the plasma reaches collisional ionization equilibrium (CIE) with $T_{\rm e} = T_{\rm z}$. The reverse case is a recombination-dominant process ($T_{\rm e} < T_{\rm z}$)---recombining plasma (RP). Previous studies on thermal X-ray spectra in SNRs have used IP to describe young SNRs and CIE plasma to describe old SNRs. This canonical scenario for SNR plasma evolution has been challenged by the recent {\it Suzaku} discoveries of RP from several mixed-morphology SNRs \citep{Rho1998}, IC\,443, W49B, G\,359.1$-$0.5, W28, W44, G\,346.3$-$0.2, and G348.5$+$0.1 \citep{Yamaguchi2009, Ozawa2009, Ohnishi2011, Sawada2012, Uchida2012, Yamauchi2013, Yamauchi2014}. As evidence of the existence of RP, the authors cite the detection of enhanced radiative recombination continua (RRCs), X-ray emissions that are made when free electrons are directly recombined with atoms in a bound state (free-bound transition). RRC is most conspicuous for the transition of free electrons to the ground state of either He- or H-like atoms. Most RRCs discovered so far are those of He-like magnesium (Mg), silicon (Si), and sulfur (S) below $\sim$4~keV, a clouded energy band in which many emission lines from abundant elements overlap on the RRC structures. Thus, the RRC signatures are only revealed as saw-teeth like residuals when fit to the IP or CIE model. On the other hand, RRC structures should be more conspicuous in the energy band 4--10 keV, because the RRC structures---those of iron (Fe) and calcium (Ca)---are more sparsely spaced, with no overlap of emission lines from the relevant elements in this energy band. However, no RRC structures of Ca and Fe have been reported---except Fe from W49B \citep{Ozawa2009}---mainly due to the limited statistics in the high-energy band. Therefore, we conducted deep observations on IC\,443, the most robust RP SNR and only the RRCs of Mg, Si, and S but no other higher Z elements have been found \citep{Yamaguchi2009}. IC\,443 (G\,189.1$+$3.0) is located on the Galactic anti-center at a distance of 1.5~kpc \citep{Welsh2003}. The remnant is associated with a dense giant molecular cloud \citep{Cornett1977} near the Gem OB1 association \citep{Humphreys1978}. Thus, IC\,443 is likely a remnant of a core-collapse supernova. Using the {\it ASCA} satellite, \citet{Kawasaki2002} found that the K-shell intensity ratios of H-like Si and S relative to He-like Si and S were significantly higher than those expected in the CIE plasma of the electron temperature determined from the bremsstrahlung continuum. They concluded that the plasma in IC\,443 was RP. \citet{Yamaguchi2009} found RRCs of Mg, Si, and S and estimated that $kT_{\rm e}$ and $kT_{\rm z}$ were, respectively, $\sim$0.6 and $\sim$1.0--1.2~keV, confirming the existence of RP. In this paper, we report further evidence of RP on the basis of the new discoveries of RRCs from Ca and Fe based on {\it Suzaku} deep observations. We discuss the characteristic features of RP. | Based on deep observations, we discover a strong He-$\alpha$ of Fe at $\sim$6.7~keV. Because the electron temperature is very low---at most $\sim$0.7~keV---it is almost impossible to emit He-$\alpha$ of Fe by collisional excitation, a dominate process to produce this line in IP or CIE plasma. Therefore, the presence of this line itself already indicates that the plasma is neither IP nor CIE. Furthermore, we discover RRCs of He-like Fe and Ca and H-like Ca, which is robust evidence for RP. The RRC of He-like Fe is the second sample after W49B, while the RRCs of He- and H-like Ca are the first discoveries in SNRs. Although the line flux of Fe He-$\alpha$ has sufficient statistics, the best-fit 1$\sigma$ error of the Fe abundance is significantly large. This situation is somewhat similar to those of Ca and Ni. This apparent ``inconsistency'' is due to the RP proper characteristic, and thus is not found in IP. The ionization rate from less ionized atoms is higher than that from more ionized atoms, and thus the number of charges in IP increases monotonically as $n_{\rm e}t$ increases. At any $n_{\rm e}t$, no large dispersion of the charge number appears. Unlike the ionization rate, the recombination rate does not strongly depend on the charge number, and thus a broad distribution of different ions is achieved in RP at a large $n_{\rm e}t$. In our case, H-like Fe can survive in a wide range of $n_{\rm e}t$, significantly varying from 10$^{11}$ to 10$^{12}$~cm$^{-3}$~s. Near $\sim$10$^{12}$~cm$^{-3}$~s, the fraction of H-like Fe decreases to almost 0\%. RRCs of He-like atoms occur due to the recombination of free electrons to the ground states of H-like ions. At low electron temperature (0.65~keV), collisional excitation rate of He-like atoms from ground to excited states is far smaller than that of the recombination process. The contribution of the cascade line, which originates from electrons captured at the excited levels of ions by free-bound transition \citep{Ozawa2009}, is significantly larger than that of the line flux that originates from collisional excitation. Therefore, the flux of He-like K-shell lines and RRC are approximately proportional to the fraction of H-like ions. Accordingly, in RP, atomic abundances and $n_{\rm e}t$ are strongly correlated---larger $n_{\rm e}t$ leads to greater abundance. We plot the confidence contour in the two-dimensional space of the relaxation parameter $n_{\rm e}t$ and the metal abundance of Fe (Figure~\ref{figure_net_Fe}). \begin{figure} \begin{center} \includegraphics{figure5.eps} \caption{ Solid lines show the error contours for the Fe abundance--$n_{\rm e}t$ space (left {\it y}-axis) and dashed lines show the abundance ratio of $Z_{\rm Ni}$/$Z_{\rm Fe}$--$n_{\rm e}t$ space (right {\it y}-axis), both in double-logarithmic scale. Confidence levels are at 1$\sigma$ (red), 2$\sigma$ (green), and 3$\sigma$ (blue). The crosses show the best-fit parameters } \label{figure_net_Fe} \end{center} \end{figure} The detection of the Ni He$\alpha$ line and a large ratio of $Z_{\rm Ni}$/$Z_{\rm Fe}$ relative to the solar abundances of $\sim$10 are also discoveries from IC\,443 and are shown in Figure~\ref{figure_net_Fe}. To check this high ratio of $Z_{\rm Ni}$/$Z_{\rm Fe}$, we compare SPEX with the other sets of atomic data concerned \citep{Masai1997,Bryans2009} and estimate the uncertainty to be 20\%. We thus conclude that the ratio of $Z_{\rm Ni}$/$Z_{\rm Fe}$ relative to the solar abundances is in the range of $\sim$8--12. Unlike the absolute value of abundance for each element, this high ratio of $Z_{\rm Ni}$/$Z_{\rm Fe}$ is valid in the error range of $n_{\rm e}t$. However, such a high ratio is not predictable from any theoretical model of spherically symmetric explosion in core-collapse supernovae \citep[e.g.,][]{Woosley1995}. Recently, a similar high ratio of Ni to Fe relative to the solar abundances $\sim$8 was obtained from core-collapse SNRs, G\,350.1$-$0.3 and G\,349.7$+$0.2, suggesting that a significant fraction of Ni is ejected from the core region of their progenitors \citep{Yasumi2014}. For SN\,2006aj, \citet{Maeda2007} proposed a model in which a large amount of $^{58}$Ni might be ejected from the core as a result of asymmetric explosion. A possibility is that IC 443 is also a remnant of an asymmetric explosion of a core-collapse supernova. | 14 | 3 | 1403.5850 |
1403 | 1403.3067_arXiv.txt | We present the first \emph{ab initio} calculations for p-shell single-$\Lambda$ hypernuclei. For the solution of the many-baryon problem, we develop two variants of the no-core shell model with explicit $\Lambda$ and $\Sigma^+,\Sigma^0,\Sigma^-$ hyperons including $\Lambda$-$\Sigma$ conversion, optionally supplemented by a similarity renormalization group transformation to accelerate model-space convergence. In addition to state-of-the-art chiral two- and three-nucleon interactions, we use leading-order chiral hyperon-nucleon interactions and a recent meson-exchange hyperon-nucleon interaction. We validate the approach for s-shell hypernuclei and apply it to p-shell hypernuclei, in particular to $\isotope[7][\Lambda]{Li}$, $\isotope[9][\Lambda]{Be}$ and $\isotope[13][\Lambda]{C}$. We show that the chiral hyperon-nucleon interactions provide ground-state and excitation energies that generally agree with experiment within the cutoff dependence. At the same time we demonstrate that hypernuclear spectroscopy provides tight constraints on the hyperon-nucleon interactions. | 14 | 3 | 1403.3067 |
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1403 | 1403.3856_arXiv.txt | We present \chandra\ \x\ point source catalogs for 9 Hickson Compact Groups (HCGs, 37 galaxies) at distances $34 - 89$~Mpc. We perform detailed \x\ point source detection and photometry, and interpret the point source population by means of simulated hardness ratios. We thus estimate \x\ luminosities (\lx) for all sources, most of which are too weak for reliable spectral fitting. For all sources, we provide catalogs with counts, count rates, power-law indices ($\Gamma$), hardness ratios, and \lx, in the full ($0.5-8.0$ keV), soft ($0.5-2.0$ keV) and hard ($2.0-8.0$ keV) bands. We use optical emission-line ratios from the literature to re-classify 24 galaxies as star-forming, accreting onto a supermassive black hole (AGNs), transition objects, or low-ionization nuclear emission regions (LINERs). Two-thirds of our galaxies have nuclear \x\ sources with \swift/UVOT counterparts. Two nuclei have \lxte~$ > 10^{42}$~\lunits, are strong multi-wavelength AGNs and follow the known \aox--\nulnu$_{({\rm near UV})}$ correlation for strong AGNs. Otherwise, most nuclei are \x\ faint, consistent with either a low-luminosity AGN or a nuclear \x\ binary population, and fall in the \lq\lq non-AGN locus\rq\rq\ in \aox--\nulnu$_{({\rm near UV})}$ space, which also hosts other, normal, galaxies. Our results suggest that HCG \x\ nuclei in high specific star formation rate spiral galaxies are likely dominated by star formation, while those with low specific star formation rates in earlier types likely harbor a weak AGN. The AGN fraction in HCG galaxies with $M_R \le -20$ and \lxte~$\ge 10^{41}$~\lunits\ is \aer{0.08}{+0.35}{-0.01}, somewhat higher than the $\sim 5\%$ fraction in galaxy clusters. | By virtue of their selection criteria, Hickson Compact Groups (HCGs) constitute a distinct class among small galaxy agglomerations. The Hickson catalog \citep{hickson1982,hickson1992} comprises 92 spectroscopically confirmed, nearby (median redshift \zmed~$=0.03$, $\sim 130$ Mpc) compact groups with three or more members with accordant redshifts (i.e., within 1000~\kmps\ of the group mean). The characteristic physical properties of CGs \citep{hickson1992} include galaxy separations of the order of a few galaxy radii (median projected separations $\sim 40h^{-1}$ kpc), low velocity dispersions (radial median $\sim 200$ \kmps) and high galaxy number densities (up to $10^8 h^2$ Mpc$^{-2}$). These conditions favor galaxy interactions, as demonstrated by the spectacular examples of HCG 92 \citep[Stephan's Quintet, e.g.][]{fedotov2011} and HCG 31 \citep{gallagher2010}. It is then natural to ask what influence this interaction-prone environment has on processes related to star-formation or accretion onto a nuclear supermassive black hole. With regards to star formation, recent work suggests that, compared to non-compact group environments, star formation is accelerated, leading to rapid exhaustion of the gas supply sustaining star forming activity. This result follows from ultraviolet and infrared star-formation estimates that show significant discontinuities in mid-infrared colors and ultraviolet+infrared specific star formation rates \citep[SSFRs,][]{johnson2007,tzanavaris2010,walker2010,walker2012}. In particular, the discontinuities indicate a bimodality between galaxies with high levels of star formation and those with little star formation. The latter have also been found to exhibit high levels of \lq\lq \hone\ deficiency\rq\rq, \defhone, as defined by \citet{verdes2001}. These authors predict an expected \hone\ mass for field galaxies of a given morphological type and compare it to the \hone\ mass of compact group galaxies, thus calculating \defhone. Taken together, the lack of galaxies with intermediate mid-infrared colors and SSFRs, as well as the high \defhone\ values are suggestive of accelerated and then abruptly truncated star formation. \begin{deluxetable*}{cccc cccc} \tablecolumns{8} \tablewidth{0pc} \tablecaption{Chandra observation log for this HCG sample \label{tab-xdata}} \tablehead{ \colhead{HCG ID} &\colhead{Obs. ID} &\colhead{Obs. Start Date} &\colhead{Detector} &\colhead{Obs. time (ks)} &\colhead{Obs. type} &\colhead{PI} &\colhead{References} \\ \colhead{(1)} &\colhead{(2)} &\colhead{(3)} &\colhead{(4)} &\colhead{(5)} &\colhead{(6)} &\colhead{(7)} &\colhead{(8)} } \startdata HCG 7 & 8171 & 2007-09-13 & ACIS-S & 19.4 & GTO & Garmire & \\ HCG 7 & 9588 & 2007-09-16 & ACIS-S & 16.9 & GTO & Garmire & \\ &&&& {\bf 36.3} &&& \citet{konstantopoulos2010} \\ HCG 16 & 923 & 2000-11-16 & ACIS-S & {\bf 12.7} & GO & Mamon & \citet{jeltema2008} \\ HCG 22 & 8172 & 2006-11-23 & ACIS-S & {\bf 32.2} & GTO & Garmire & \citet{desjardins2013} \\ HCG 31 & 9405 & 2007-11-15 & ACIS-S & {\bf 36.0} & GO & Gallagher & \citet{smith2012} \\ HCG 42 & 3215 & 2002-03-26 & ACIS-S & {\bf 32.1} & GO & Ponman & \citet{jeltema2008} \\ HCG 59 & 9406 & 2008-04-12 & ACIS-S & {\bf 38.9} & GO & Gallagher & \citet{desjardins2013} \\ HCG 62 & 921 & 2000-01-25 & ACIS-S & 49.1 & GO & Vrtilek & \\ HCG 62 & 10462 & 2009-03-02 & ACIS-S & 68.0 & GO & Rafferty & \\ HCG 62 & 10874 & 2009-03-03 & ACIS-S & 52.0 & GO & Rafferty & \\ &&&& {\bf 169.2} &&& \citet{jeltema2008}\\ HCG 90 & 905 & 2000-07-02 & ACIS-I & {\bf 50.2} & GO & Bothun & \citet{jeltema2008} \\ HCG 92 & 7924 & 2007-08-17 & ACIS-S & 94.4 & GO & Vrtilek &\\ HCG 92 & 789 & 2000-07-09 & ACIS-S & 20.0 & GO & Trinchieri &\\ &&&& {\bf 114.4} && & \citet{osullivan2009} \enddata \tablecomments{Columns are: (1) HCG group name; (2) observation ID; (3) start date of observation; (4) detector; (5) exposure time; (6) observation type (Guarranteed Time observing or General Observer proposal; (7) principal investigator; (8) references (first publication using these data). Total exposure times for each group appear in bold. } \end{deluxetable*} The importance of accretion onto a nuclear supermassive black hole (SMBH) in compact groups (\lq\lq AGN\rq\rq\footnote{In line with common usage in the literature we shall use the acronym \lq\lq AGN\rq\rq\ (active galactic nucleus) to refer to accretion onto a nuclear supermassive black hole. Strictly this is incorrect as nuclear activity can also be due to star formation.}) has not been thoroughly investigated and is not well established. In galaxy clusters \citet{dressler1985} found fewer AGNs compared to the field \citep[but see also][and below]{martini2006}. Compared to clusters, compact groups of galaxies have lower velocity dispersions making prolonged close interactions more likely. It is thus possible that the level of AGN activity is different. On the theoretical and computational side, simulation work \citep[e.g.][]{hopkins2010} suggests that major galaxy mergers are a leading mechanism that can trigger inflow of rotationally supported gas to feed a central SMBH. Note though that this would also provide fuel for intense star formation and could trigger nuclear starbursts \citep[e.g.][]{mihos1996}. Other feeding mechanisms include supernova winds, minor interactions, and disk instabilities. Several observational surveys have provided insight on the connection between AGNs and galaxy interactions. For instance \citet{kartaltepe2010} find that AGNs are common in Ultraluminous and Hyperluminous Infrared Galaxies (ULIRGs and HyLIRGs), which are known to result from major mergers. In addition, the AGN fraction in this population increases with infrared luminosity. Recently, \citet{silverman2011} find increased AGN activity in pairs compared to isolated galaxies. On the other hand, several authors find minor interactions and secular evolution to be most important in triggering AGN activity \citep[e.g.][]{grogin2005,georgakakis2009,cisternas2011,deng2013}. In the optical regime, \citet{coziol1998a,coziol1998b,coziol2004} used emission-line ratios in several samples (up to 91 galaxies in 27 compact groups) to determine the type of nuclear activity in compact group galaxies, consistently finding that strong and low-luminosity (\lha~$\lesssim 10^{39}$~\lunits) AGNs (LLAGNs) each make up no more than $\sim 10$\%\ of the total CG populations \citep[see][Table 3]{coziol2004}. Depending on the specific sample, star-forming galaxies represent a fraction up to $\sim 34$\%\ of the population, with the remaining galaxies showing no emission lines. Both LLAGNs and AGNs are found mainly in optically luminous early type galaxies with little on-going star formation that are in the centers of evolved groups. This finding was interpreted to indicate that such group cores are old, collapsed systems where star formation activity has ceased. According to this interpretation, high central densities of group cores induced gravitational interactions, which accelerated star formation, rapidly consuming all of the available fuel. It is important to note that the fractions for \lq\lq LLAGNs\rq\rq\ presented by these authors also include low-ionization nuclear emission regions (LINERs), the nature of which is still a matter of debate. LINERs are characterized by high ratios of narrow optical low ionization oxygen emission lines \citep{heckman1980} and are found in about half of all nearby galaxies \citep{ho1997}. Candidate power sources for LINERs include (1) weak AGNs \citep[e.g.][]{halpern1983,ferland1983}, (2) hot stars \citep[e.g.][]{terlevich1985,filippenko1992,shields1992}, and (3) shocks \citep[e.g.][]{heckman1980,dopita1996}. Although weak AGNs have been found in the majority ($\sim 75$\%) of LINERs \citep[e.g.][]{barth1998,ho2001,filho2004,nagar2005,maoz2005,flohic2006,gonzalez-martin2009}, they cannot account for the total LINER emission in the majority of cases \citep{eracleous2010a}. In fact, for most LINERs \citet{eracleous2010a} show that there is an energy deficit problem: Star formation and AGN activity are not able to provide a sufficient number of ionizing photons to account for the observed emission lines. In the most recent optical study \citet[][hereafter M10]{martinez2010} compiled a large spectroscopic sample of 280 galaxies in 64 HCGs and used emission-line ratios to classify the type of nuclear activity, providing an estimate for the AGN fraction in HCGs. They classified 23\%\ of galaxies as AGNs, % 10\%\ as transition objects (TO), % and 14\%\ as star forming (SF), % with the remainder of the galaxies showing no emission lines. According to this study, although AGNs appear to be the most numerous emission-line galaxy class in CGs, they have characteristically low \ha\ luminosities (median \ten{7.1}{39}~\lunits) and virtually no broad emission lines, suggestive of LLAGNs. However, these authors use a restricted set of line ratios that precludes distinguishing between LINERs and AGN. In this paper we use the \citet{kewley2006} method to reclassify the galaxies of M10. This allows us to also identify LINER systems. To stress that this is an optically based classification, we use the designations optAGN, optTO, optSF, optLINER. Work in different wavelength regimes can provide complementary insight into these questions. \citet{gallagher2008} used $1-24$\micron\ 2MASS+\spitzer\ nuclear data to probe the nuclear activity in 46 galaxies from 12 nearby HCGs. They found that the spectral index, \airac, of a power law fit to the $4.5-8.0$\micron\ IRAC data cleanly separates MIR-active from MIR-inactive HCG nuclei. Unfortunately, the exact origin of activity (whether AGN or star-formation) cannot be deduced by this method. In particular, these authors show that hot dust emission can be responsible for their results, and this can be due either to hard ionizing AGN continua or AGB populations in star forming galaxies. On the other hand, \citet{roche1991} have shown that MIR-inactivity (\airac~$>0$) is associated with {\it low}-luminosity AGN activity. Due to the high-energy emission generated by supermassive black hole accretion, by far the best direct diagnostic for strong AGN activity is nuclear \x\ emission. Compared to the optical, the \x\ regime offers the advantage that the nuclear emission is not diluted by starlight from the host galaxy, while dust obscuration is very significantly mitigated due to the higher, penetrating power of \x\ radiation. Unfortunately, this simple picture is complicated by the combined effect of two factors. First, \x\ starlight sometimes can actually dilute AGN emission. This is because \x\ binary (XRB) populations in circumnuclear star clusters also emit in the \x\ regime, although {\it individual} XRBs typically have lower luminosities than strong AGNs. Second, as the name implies, LLAGNs emit at low \x\ luminosities. Adopting a fiducial threshold of \lxte~$=10^{41-42}$~\lunits, it is only at higher \x\ luminosities that nuclear \x\ emission can be attributed to an AGN with high probability. Thus the situation becomes increasingly ambiguous at progressively fainter luminosities, making it challenging to distinguish between \x\ emission due to unresolved populations of circumnuclear XRBs and that of LLAGNs. In this regime high angular resolution becomes critical for distinguishing nuclear from circumnuclear emission. Although earlier studies did detect \x\ emission in HCGs, they were hampered by poor angular resolution and the lack of hard \x\ sensitivity, making it difficult to disentangle the contributions from point source (nuclear or extra-nuclear) and diffuse emission, and essentially concentrated on studying the diffuse component. Using \rosat\ data, \citet{ponman1996} detected a diffuse IGM in $\sim 75$\%\ of a large HCG sample, while \citet{mulchaey2003}, using a low-redshift sample of 109 groups that included poor, compact as well as rich, non-compact systems found diffuse, extended \x\ emission in 61 groups (56\%). In an effort to understand the relevance of ram-pressure stripping and strangulation due to a hot IGM in the most \hone-deficient HCGs, \citet{rasmussen2008} also examined the level of nuclear activity in a sample of 8 HCGs, finding no significant enhancement. However, they do not carry out a detailed high angular resolution study to provide more specific results on the nature of nuclear activity in their systems. The level of AGN activity in galaxy {\it clusters} has already been systematically investigated in the \x\ regime, leading to differing conclusions \citep[e.g. see][for a review]{ehlert2013}. Using a multi-wavelength approach that includes emission lines, \x\ spectral properties and \x\ to visible-wavelength flux ratios in rich clusters, \citet{martini2006} find that $\sim 5$\% of cluster galaxies more luminous than $M_R = -20$ host AGNs with \lxte~$ > 10^{41}$~\lunits. They notably also find a discrepancy between the AGN fraction determined from optical spectroscopy and a higher fraction suggested by \x\ luminosities. Interestingly, \citet{shen2007} compare the environments of poor groups and clusters using a combined optical and \x\ approach. They conclude that poor groups host AGNs that are in an optically dominant phase, whereas those in clusters are dominant in the \x s, leading to the findings of \citet{martini2006}. In compact groups there has been to-date no systematic study of nuclear \x\ emission. In this paper we take advantage of the superb angular resolution of the \chandra\ \x\ observatory to carry out detailed point source detection in a sample of 9 compact groups (37 galaxies). This paper has two main goals: First, we make available full \x\ source catalogs based on the \chandra\ observations in 9 compact group fields with detailed information on counts, fluxes, luminosities and hardness ratios. Second, we focus on point sources located in HCG galaxy nuclei. Using \chandra\ and \swift/Ultra-Violet and Optical Telescope \citep[\uvot;][]{2005SSRv..120...95R} data, we combine \x\ and ultraviolet (UV) nuclear photometry, and compare with radio and optical diagnostics to assess the nature of nuclear activity in compact group galaxies. In a separate paper, we discuss the diffuse \x\ emission in the same sample of compact groups \citep{desjardins2013}. Some of the \chandra\ data have first been presented previously in a different context. We give appropriate references in \tr{tab-xdata}. The structure of the paper is as follows: Section 2 introduces our sample. Section 3 discusses X-ray data and analysis and point source detections. Section 4 presents UV nuclear data and analysis. Section 5 presents multiwavelength analyses, including new optical emission-line ratio classifications, radio data and a combined \x-UV analysis. Section 6 presents estimates on the AGN fraction in HCGs and Section 7 discusses our findings. We conclude with a summary in Section 8. | This paper presents the first compilation of \x\ detected point sources in the fields of 9 HCGs, for which we provide an extensive compilation of source characteristics. We have used multi-wavelength diagnostics (\x, UV, optical, and radio) to assess the levels of AGN, SF and LINER activity in the compact group environment. Our main results are the following: \begin{enumerate} \item In 60\%\ of 37 galaxies we detect single, nuclear \x\ sources that have nuclear UV counterparts. We detect no nuclear \x\ emission for 27\%\ of our galaxies. The rest of the systems have more uncertain \x\ nuclear detections. \item Out of the 22 galaxies for which emission line ratios are available in the literature, we classify a clear plurality (45.5\%) asn optSF. Our criteria allow us, for the first time, to also classify five systems as LINERs, although four of these are mixed (LINER/AGN, TO/LINER). Thus any LINER activity is associated with a minority (22.5\%) of systems. Only three nuclei (13.6\%) are classified optAGN. \item Only three systems (HCG 16~B, 90~A, 92~C) are candidates for hosting an \x\ strong AGN (\lxte~$\ge 10^{41}$~\lunits). \item When several criteria are taken into account (optical spectroscopic classification, excess radio emission, \x\ luminosity, location in \aox--\nulnu$_{2600}$ parameter space) only two HCG nuclei (90~A, 92~C) fulfill several criteria and are classified as strong, unambiguous AGNs. \item In \aox--\nulnu$_{2600}$ space, HCG nuclei occupy a region which is distinct from that occupied by strong AGNs and largely overlaps with that occupied by other, nearby star-forming galaxies not known to harbor AGNs (\fr{fig:aox_nulnu}). The only exceptions are the two strong AGNs which do fall in the AGN region. We thus tentatively make the prediction that HCG nuclei without optical nuclear-type classifications are dominated by star formation (if they have late-type morphologies) or may harbor {\it low} luminosity AGNs (especially if they have early-type morphologies). \item \aox\ anticorrelates with galaxy-wide SSFR and spiral morphology so that the star formation contribution is strongest in highest SSFR and later type morphology galaxies\fr{fig:aox3}. The detected anticorrelation (correlation) with SFR (\mstar) is weaker. \item Using the same criterion used in galaxy clusters \citep{martini2006}, the AGN fraction of HCG galaxies more luminous both than $M_R = -20$ and \lxte~$=10^{41}$~\lunits\ is \aer{0.08}{+0.35}{-0.01}, which is close but higher to that in clusters. \end{enumerate} Our general conclusion is that overall the CG environment has a mitigating effect on the level of AGN activity but not AGN numbers. With future expanded \x\ and UV samples as well as deeper observations we will be able to better assess the nature and statistics of HCG nuclei. The comparison with galaxy clusters suggests that environment plays a key role for the overall level of AGN activity. In this respect, it is imperative to carry out detailed comparisons with samples from other group, cluster and field environments. | 14 | 3 | 1403.3856 |
1403 | 1403.1581_arXiv.txt | We have measured the line-of-sight velocity distribution from integrated stellar light at two points in the outer halo of M87 (NGC~4486), the second-rank galaxy in the Virgo Cluster. The data were taken at R~=~480\arcsec\ ($\sim 41.5$~kpc) and R~=~526\arcsec\ ($\sim 45.5$~kpc) along the SE major axis. The second moment for a non-parametric estimate of the full velocity distribution is $420 \pm 23$~\kms\ and $577 \pm 35$~\kms\ respectively. There is intriguing evidence in the velocity profiles for two kinematically distinct stellar components at the position of our pointing. Under this assumption we employ a two-Gaussian decomposition and find the primary Gaussian having rest velocities equal to M87 (consistent with zero rotation) and second moments of $383 \pm 32$~\kms\ and $446 \pm 43$~\kms\ respectively. The asymmetry seen in the velocity profiles suggests that the stellar halo of M87 is not in a relaxed state and confuses a clean dynamical interpretation. That said, either measurement (full or two component model) shows a rising velocity dispersion at large radii, consistent with previous integrated light measurements, yet significantly higher than globular cluster measurements at comparable radial positions. These integrated light measurements at large radii, and the stark contrast they make to the measurements of other kinematic tracers, highlight the rich kinematic complexity of environments like the center of the Virgo Cluster and the need for caution when interpreting kinematic measurements from various dynamical tracers. | \label{sec:intro} The assembly of stellar mass in elliptical galaxies has been the subject of significant investigation in recent years. With the discovery of a massive and old elliptical galaxy population at high redshift \citep[e.g.,][]{cim02a,cim04,mcc04a} we have been forced to revisit the mechanisms of massive galaxy growth. The puzzle is multifaceted. A central question is how such massive galaxies can exist in the early universe when they are generally considered the end-products of hierarchical assembly \citep[e.g.,][]{del07a}. Moreover, many of the high redshift quiescent galaxies appear to be particularly compact in comparison to local elliptical galaxies \citep[][]{dad05,tru06,tof07,van08b,cap09,cas10,wei11}, with half-light radii of $\sim 2$~kpc and no evidence for extended stellar halos \citep[e.g.,][]{szo12}. Further observations indicate that the very central regions (e.g., R $\le 1$~kpc) of z~$\approx 0$ galaxies are not physically denser than the high redshift population \citep[][although see Poggianti et al. 2013]{hop09,van10,tir11,van13,dul13} suggesting elliptical galaxy growth occurs predominantly at large radii \citep[e.g.,][]{van10}. While there remains some debate about the degree of mass evolution in the cores of massive elliptical galaxies, the growth of mass at large radii (e.g., R $\ge 5$~kpc) is well established observationally \citep[][]{bui08,van08c,van10,new12,van13,pat13}. High spatial resolution simulations from cosmological initial conditions can recreate this mass growth at large radii, commonly through a growth history dominated by minor mergers \citep[][]{naa07,naa09,ose10,ose12}. However, both the degree of growth and how the mass is assembled over time remain poorly constrained observationally. In the case of very massive ellipticals and brightest cluster galaxies (BCG), the case is even more confounding as the BCGs at high redshift have been found to closely match the luminosity \citep[e.g.,][]{sto10}, mass \citep[][]{whi08,col09} and scale \citep[e.g.,][]{sto11} of BCGs in the local universe in certain work, but show significant growth in others \citep[][]{ber09,val10,asc11}. If we are to gain a complete understanding of how BCGs and other massive galaxies assemble their mass, we will need dynamical observations at large radii where the growth is expected to occur. To this end we have made a measurement of the velocity dispersion of the integrated starlight in M87, the second-rank galaxy in the Virgo Cluster, using the Mitchell Spectrograph (formally VIRUS-P). This work is driven in part by the results of \citet{doh09} and \citet[][hereafter S11]{str11} who find a declining velocity dispersion profile from PNe and GC data respectively. Contrasting this are the results of \citet[][hereafter MGA11]{mur11} who found a rising velocity dispersion with radius for the stars. As certain systems have exhibited good agreement between different tracers of mass \citep[e.g.,][]{coc09} and others show disagreement \citep{rom01,chu10,she10,ric11}, we set out to make a measurement of the stellar velocity dispersion of M87 at a large radial distance in order to directly compare to the results of \citet{doh09} and S11. This approach was used by \citet{wei09} to good effect on the massive local ellipticals NGC~3379 and NGC~821. The traditional estimate of M87's half-light radius (\re) from RC3 \citep{dev91} is $\sim 95$\arcsec. This estimate of the \re\ of M87 puts our measurements at $\sim 5.0$~\re\ and $\sim 5.5$~\re. However, \citet{kor09} measure an \re\ for M87 that is more than a factor of 7 times larger than the canonical value. A consideration of the deep photometry of \citet{mih05} and \citet{jan10} makes it clear that the definition of the half-light radius in the centers of clusters, particularly ones as unrelaxed as Virgo, is perhaps ill-defined. We therefore elect to not use \re\ in this work and will cite physical and/or on-sky scales where appropriate. To remain consistent with our previous papers on M87 \citep[][]{geb09,geb11,mur11} we assume a distance to M87 of 17.9~Mpc. This distance corresponds to a scale of 86.5~pc~arcsec$^{-1}$. All references to other values in the literature are scaled accordingly. The paper outline is as follows. In \S~\ref{sec:data} we describe the observations and data collection. \S~\ref{sec:reductions} outlines the data reduction steps. Our results are found in \S~\ref{sec:results} where we also place our results in context with several velocity dispersion values from the literature. A discussion of the implications of these results in the context of both the assembly of mass in massive galaxies and comparisons of the different tracers of mass is found in \S~\ref{sec:discussion}. We summarize our results in \S~\ref{sec:summary}. \vspace{1cm} | The asymmetry of the stellar velocity profiles indicates complex stellar dynamics at the position of our pointing. In an analysis of the stellar dynamics we are confronted with two coupled issues. First is the issue of an accurate measurement of the second moment of the stellar component of M87. This measure is complicated by the strong asymmetry in the LOSVD. We have given the non-parametric measurement of the second moment as the most direct interpretation of the velocity profile. The parametric, Gauss-Hermite measure of the velocity profile also returns very high values for the second moment. This brings us to the second issue. That is, the intriguing possibility that the strong asymmetry in the velocity profile stems from a cooler, second component of stars superimposed along the line-of-sight and offset to a positive velocity of $\sim 1000$~\kms\ from the rest velocity of M87. The evidence for a second component is explored in \S~\ref{sec:formation} and \S~\ref{sec:conflict}. For now we simply note that if there is a second component, fitting our velocity profiles with two Gaussians yields a more representative measure of M87's stellar halo temperature. Yet even this most conservative measure of the stellar velocity dispersion ($383 \pm 32$~\kms\ and $446 \pm 43$~\kms) is still well above both the gradually declining GC velocity dispersion values of S11, and the PNe measurements of \citet{doh09}. What can we make of these conflicting results between the stars, GCs and PNe kinematics? The discrepancies between our stellar velocity dispersion values and both the GC kinematics from S11 and PNe measurements from Doherty et al. are intriguing, but not entirely surprising; there is extensive evidence indicating that the center of the Virgo Cluster is still in active formation \citep{tul84,bin87,bin93,wei97,mih05,doh09,rud10,kra11,rom12}. In considering our results in this light, we believe our pointing falls on a dynamically hot and complex region of M87 and thus reduces the tension between these potentially disparate data sets. We outline some of the relevant work done on the dynamics of the center of M87 in \S~\ref{sec:formation} and explore the tension between these and previous results in \S~\ref{sec:conflict}. Then in \S \ref{sec:rising} we explore the evidence for a rising velocity dispersion and massive dark matter halo. We then compare the connection between M87 and several BCGs that exhibit a similar rising velocity dispersion profile. \subsection{A System in Formation}\label{sec:formation} In the deep photometry of \citet{mih05} there is evidence of elongated and disturbed isophotes towards the SE major axis. Moreover, superimposed on the galaxy light of M87 is extensive intracluster light (ICL) \citep{mih05,jan10,rud10}, formed from stripped stars bound to the Virgo Cluster rather than a specific galaxy. The ICL has been studied extensively over the past decade \citep[for a nice overview, see][]{arn10} and is predicted to be ubiquitous in galaxy clusters \citep[e.g.,][]{wil04}. In \citet{wei97} they find a stellar stream extending to nearly 100~kpc along the SE major axis, directly across our field (see Figure 1 in that work) and speculate this stellar debris comes from the recent accretion of a spheroidal galaxy. With a short dynamical time for the extended material (t~$\le 5 \times 10^8$ years) the authors argue we have either caught M87 during a special time during its formation, or that these events are common in the buildup of the outer halos of massive ellipticals. As recent theoretical work on the mass assembly of ellipticals points to minor mergers as a primary mode of mass assembly \citep{naa09,ose10,joh12}, and that the bulk of the ICL is built by the stripping of stars during the formation of the BCG \citep[e.g.,][]{mur07} it appears likely these events are common. Further evidence of the continued assembly of M87 comes from the GC kinematical analysis of \citet{rom12} where the phase-space substructure of the M87 GC population reveals the existence of at least 2 components. In their work (see their Figure 1) a shell-like structure is discovered. Their formation simulations seem to indicate the accretion of a $\sim 0.5$~L$^*$ elliptical progenitor with an observable lifetime of $\sim 1$~Gyr. As the authors point out, there is difficulty in reconciling the large number of GCs in the shell component with the relatively cold velocity dispersion they measure. One plausible scenario they suggest is that the accretion of an $\sim$~L$^*$ galaxy (i.e. a $\sim 1:10$ merger) combined with the stripping of GCs from several satellite compact dwarf galaxies could lead to both a large number and simultaneously cool GC component. Extensive work on the dynamics of the Virgo Cluster as a whole,\footnote{See the introduction in \citet{fou01} for a good overview of the studies of Virgo galaxies.} particularly the pioneering work of Binggeli and collaborators \citep{bin85,bin87,bin93}, has shown clear evidence that the Virgo Cluster is young and unrelaxed. The Virgo cluster is not unique in this regard and signs of continued assembly can be seen in other galaxy clusters such as the Coma Cluster \citep[e.g.,][]{ger07}. Intriguingly, the distribution of Virgo galaxy velocities shows a pronounced asymmetry to positive velocities, with a peak offset by $\sim 1000$~\kms\ from the Virgo recession velocity \citep{bin93,fou01,mei07}. Although a much more comprehensive dynamical analysis is needed to say anything definitive, the idea that one of these galaxies was stripped during a close passage to M87, depositing a relatively cool stellar component at a velocity offset of $\sim 1000$~\kms, is one possible source of the velocity profile asymmetry we see. \subsection{Conflicting Results?}\label{sec:conflict} Neither the GC measurements of S11, nor the PNe measurements of \citet{arn94} and \citet{doh09}, are in line with our stellar velocity dispersion measurements. Moreover, the GC measurements of \citet{han01}, as rebinned for the work of MGA11 and shown in Figure \ref{fig:disp}, are significantly higher than 3 of our 4 stellar measurements. We now turn to each set of conflicting results in search of some relief to this tension in the observations. As the discrepancy between the PNe measurements and the stars is the most dramatic, we will begin there. Due to the evidence of continued assembly of the Virgo Cluster, we believe the comparison to the PNe data from \citet{doh09} is simply ill-advised; their data comes from the opposite side of the galaxy as our stellar kinematics (see their Figure 7) and is likely dynamically unrelated to the stars in our field. Also of note is the difference in radial position between the PNe and stellar measurements. The Doherty et al. pointings are centered at 13.7\arcmin\ and 32.9\arcmin\ compared to our data points at 8.0\arcmin\ and 8.8\arcmin. When seen in this light, agreement between the PNe and stellar measurements would be more surprising than not. Next we explore the conflicting results between the GC measurements of S11 and \citet{han01}. As stated earlier, in order to match the dynamical modeling bins of MGA11 we rebinned the individual Hanes et al. GC velocities. These values are plotted in Figure \ref{fig:disp} and are slightly different than what is found in the dynamical analysis of \citet{cot01}. As discussed in S11, the Hanes et al. values were found to contain a handful of ``catastrophic outliers''. Once these outliers are removed, and their new GC data included, their measured dispersion drops substantially. This relieves the tension between the GC measurements, but still does not explain why our stellar velocity dispersion values are significantly higher than the S11 GC values. Unlike the PNe values, the spatial agreement between the S11 GCs and our field is good; we have 4 of their GCs within 200\arcsec\ from the center of our field. These coincident GCs are very cold and bare no resemblance to the dynamics we see in the stars.\footnote{The GCs recession velocities (relative to the rest velocity of M87) within 200\arcsec\ of our field from S11 are as follows. S87: $83\pm103$~\kms, S93: $-12\pm48$~\kms, S170: $-22\pm106$~\kms, S270: $86\pm43$~\kms.} However, this apparent conflict between the stars and GC kinematics of S11 gets some relief by taking a step back and considering the GC population of M87 as a whole. There are a couple of indications that the 8\arcmin~$\le$~R~$\le$~10\arcmin\ region of M87 is dynamically hotter and more complicated than other regions of the galaxy. Figure 23 in S11 shows the GC LOSVDs for the 4 subpopulations they define in their sample. In their inner radial bin (R~$\le$~10\arcmin) both the blue GCs, bright GCs (i$_0 < 20$) and UCDs show distinctive wings to positive velocity, similar to the LOSVDs of the stars in this work. The wings to positive velocity are also seen in the GC velocity profiles of \citet[][Figure 2]{rom01}. Although the agreement between the degree of offset in the velocity profile asymmetry seen in the S11 GCs and our stars is not perfect, it suggests a possible link between the mechanisms responsible for assembling the blue GCs, bright GCs, UCDs, and stars at this radius. This connection between the stars and blue GCs is also seen in the spatial distribution and chemical analysis of \citet{for12}. In considering the possibility of such a link more closely, Figure 21 in S11 shows the position angle (PA), rotation velocity, velocity dispersion and velocity kurtosis of their 4 subpopulations of GCs. There are two points we make here. First, the blue GC population (left-most column) shows a distinct change in PA just beyond R~$> 10$\arcmin. This change in PA corresponds to a spike in rotational velocity while the velocity dispersion drops at $\sim 10$\arcmin\ from $\sim 370$~\kms\ (in rough agreement with our two-Gaussian measurement) to $\sim 300$~\kms. Second, of keen interest to us are the UCDs and bright GCs in their sample (right-most column in their Figure 21) which show a strong rise in velocity dispersion to well above 400~\kms\ at $\sim 8$\arcmin. At the same radius the bright GCs reach nearly 500~\kms, and both the UCDs and bright GCs show a noticeable spike in rotation velocity at this position. This spike in velocity dispersion is also shown in the left-hand plot of their Figure 20 where \emph{both} their faint and bright GCs become quite hot, right at R~$\approx 8$\arcmin. These signatures of a change in the dynamical nature of M87 were noted by S11 where they suggest that the 4 dwarf ellipticals found between 7\arcmin\ and 9\arcmin\ are ``stirring the pot''. This leads us to explore the UCD population as not only the cause of the high stellar velocity dispersion, but also the source of the velocity profile asymmetry. M87 has a radial velocity of 1307~\kms\ \citep[][see also Mei et al. 2007 and Makarov et al. 2011]{huc12}. The UCDs nearest our field are NGC~4486a and IC~3443 (see Figure \ref{fig:mihos}). The radial velocity of NGC~4486a is 757~\kms\ \citep{pru11} and therefore can not be the source of the positive velocity asymmetry. A radial velocity of 2272~\kms\ was reported for IC~3443 in \citet{dev91} which was in good agreement with the earlier work of \citet[][]{eas78} who measured a value of 2254~\kms. At these values, the halo of IC~3443 would appear to be a candidate for the source of the $\sim 1000$~\kms\ wing to positive velocity seen in our velocity profiles. However, more recent radial velocity measurements for IC~3443 return lower values, typically under 2000~\kms\ \citep{bot88,van00,gav04}, with 1785~\kms\ being the currently accepted value \citep{ade06,rin08,aih11}. At this radial velocity, stars from the halo of IC~3443 can not explain the velocity profile wing. In order to determine if there is any photometric evidence for a second population of stars within our field we have inspected the deep photometry of \citet{mih05} to look for any variation beyond the smooth decline in the surface brightness of M87. Despite the aforementioned evidence for disturbed isophotes along the SE edge of the major axis \citep{wei97,mih05}, we see no distinct changes across our field. We also find no significant deviations from a smooth decline in the surface brightness with radius from an inspection of the flux in our fibers shown in Figure \ref{fig:fibflux}. However, despite not finding evidence for a kinematic disturbance in the photometry, \citet{rom12} point out in their discussion that M87 is a favored target for its proximity and ``general lack of obvious dynamical disturbance [sic]''. Yet dynamical studies have now revealed rich phase-space complexity in not only M87 but several other massive ellipticals that is not apparent in the photometry alone. As our new data points exhibit kinematic complexity not seen in the velocity profiles from MGA11 it leads us to wonder whether the outshirts of the stellar halo of M87 hold a wealth of clues to its formation. Clearly, further observations of the stellar halo of M87 are warranted. \subsection{A Rising Stellar Velocity Dispersion}\label{sec:rising} We have quoted 3 different estimates of the stellar velocity dispersion of M87 for our R1 and R2 fields: a non-parametric measure of $420 \pm 23$~\kms\ and $577 \pm 35$~\kms, a Gauss-Hermite parameterization of $456 \pm 37$~\kms\ and $604 \pm 45$~\kms, and a two-Gaussian parameterization of $383 \pm 32$~\kms\ and $446 \pm 43$~\kms. We advocate the non-parametric measure as the most straightforward interpretation and also the most relevant for dynamical modeling. However, due to the intriguing, albeit tentative, evidence for the existence of a second stellar component, we remain somewhat agnostic about which fit better represents the true velocity dispersion of the stars in the M87 halo. That said, \emph{even the lowest estimate of the dynamical temperature of the stars at R~$\approx$~500\arcsec\ shows a clear increase to above 400~km~s$^{-1}$.} What can we make of this continued increase in stellar velocity dispersion out to nearly 45~kpc in M87, and how does it inform our picture of galaxy structure and formation? \subsubsection{The Dark Matter Halo of M87}\label{sec:dmhalo} In terms of galaxy structure, if the rising velocity dispersion profile reflects the gravitational potential of the galaxy, then the presence of a massive dark matter (DM) halo can explain such a rise. Indeed, a very massive DM halo has been detected in M87 by several groups employing a variety of methods \citep{fab83,huc87,mou87,mer93,rom01,mat02,das10,mur11}. Yet clearly the velocity dispersions shown in Figure \ref{fig:disp} can not all be reflections of the gravitational potential of M87. One relevant point here is that in the case of both GC and PNe measurements one is working with individual data points. While the GC population of this and other massive early-type galaxies can reach several thousand \citep[e.g.,][]{mcl99}, measurements of their kinematics are typically done with a few hundred GCs. In the work of \citet{doh09}, their PNe measurement of velocity dispersion comes from 12 PNe for the 247~\kms\ measurement and 9 for the lower 139~\kms\ value. S11 are in a far better position, having velocity measurements of over 700 GCs. With the kinematic complexity of M87, sampling sub-populations that are not reflective of the gravitational potential of the galaxy is a possibility. With integrated starlight we avoid these statistical challenges. Yet if the 8\arcmin~$\le$~R~$\le$~10\arcmin\ region is dynamically hotter, and with our much smaller field-of-view than covered by either the GC or PNe populations, we find ourselves in a similar position of not fully sampling the phase space. Certainly more data and a more complete dynamical analysis is in order, but for now we proceed to interpret these results with the understanding that more stellar data is required. Further support for the presence of a massive DM halo comes from the X-ray gas mass estimates of \citet[][]{mat02}. Their mass estimates align well the with \citet{coh97} GC velocity dispersion profiles (see their Figure 21) which were compiled into the Hanes et al. values. As larger radial data points provide greater leverage on the total enclosed mass of a galaxy, the GC values of Hanes et al. play a significant role in constraining the DM halo measured in MGA11. In that work we plot a comparison of our best enclosed-mass profile to a variety of literature values (Figure 11 and Table 4 of MGA11). The MGA11 enclosed mass estimate for either an NFW or cored-logarithmic DM halo are generally in good agreement with the mass estimates from the literature, particularly at large radii \citep[e.g.,][]{fab83}. However, S11 conduct a similar comparison (see their Figure 16) and find a substantially less massive DM halo for M87. This lower mass is certainly driven by their lower velocity dispersion. As those authors point out, comparisons between data sets from various groups is challenging, and perhaps attempting to align the various mass tracer populations in M87 is ill-advised; the kinematics of M87, and the center of the Virgo Cluster, are complicated and there is no a~priori reason that the GC and stellar kinematics must go in lock step as different formation pathways will leave different kinematic signatures. For example, the accretion of smaller satellite galaxies with cooler kinematic components can lead to a lower velocity dispersion in the GCs than those formed in~situ, as pointed out in \citet[][]{rom12}. These complications necessitate both a more complete set of observations of the kinematic components and comprehensive dynamical modeling in order to get the entire picture of M87's formation history. Yet with the good agreement between the gas kinematics of \citet[][]{mat02} and the stellar kinematics of MGA11 and those presented here, it appears M87's DM halo is the dominant mass component by 30~kpc. \subsubsection{M87 as the BCG of Virgo?}\label{bcgs} How does M87 compare to other central galaxies? Although M87 is not technically the BCG of Virgo\footnote{M49 is slightly more luminious and thus the ``brightest'' galaxy in the Virgo Cluster \citep{kor09}.} it does occupy a central location in terms of cluster mass \citep[e.g.,][]{boh94}. We therefore compare M87 to other BCGs and find that rising stellar velocity dispersions are not uncommon. Early work by \citet{dre79} on the BCG in the Abell~2029 galaxy cluster (IC~1101) found a stellar velocity dispersion that increases to over 500~\kms\ at 100~kpc. \citet{kel02} find a similar result for NGC~6166, the BCG in Abell~2199 \citep[see also][]{car99}, where the stellar velocity dispersion gradually rises to $\sim 660$~\kms\ by 60~kpc. Also, from stellar velocity dispersion measurements and strong lensing constraints, \citet{new11} find the BCG in Abell~383 to exhibit a steeply rising dispersion profile that climbs from $\sim 270$~\kms\ at the center of the galaxy to $\sim 500$~\kms\ by $\sim 22$~kpc. And the extensive work of \citet[][]{lou08} on 41 BCGs with long-slit data find a significant fraction of their sample exhibit flat to rising velocity dispersion profiles.\footnote{In their paper Loubser et al. claim to find 5 galaxies with rising stellar velocity dispersions. However, if we allow for even a gradual rise in velocity dispersion and inspect their figures, we find this number increases to $\sim 14$.} Obviously, rising velocity dispersions in other galaxies do not give support for a rising velocity dispersion in M87. We simply want to highlight that as observations improve, and dynamical measurements at ever-larger radii become possible, we have been finding rising velocity dispersion profiles in BCGs in greater abundance. This leads us to a final comparison between M87 and another BCG. In our non-parametric measure of the second moment, we see a rise of $\sim 150$~\kms\ over 4.3~kpc. This sharp rise in velocity dispersion over a relatively short radial distance is quite striking, but not unprecedented. In \citet{ven10} the stellar kinematics of the BCG NGC~3311 are found to rise very rapidly from $\sim 150$~\kms\ at the center to $\sim 450$~\kms\ at R~$\approx~13$~kpc. Over a very similar physical distance (from R~$\approx 8$~kpc to R~$\approx 13$~kpc, a difference of $\sim 5$~kpc) the velocity dispersion rises from around 280~\kms\ to above 450~\kms, a steeper rise than we see for M87. \citet{ric11} remeasure the stellar velocity dispersion of NGC~3311 and use GC velocities to constrain the DM halo, and although they report considerably lower values for the overall stellar velocity dispersion than \citet{ven10}, they see the same steep rise over a similar radial distance. Also noteworthy in NGC~3311 is the difference between the stars and GCs; the GC velocity dispersion at R~$\approx~15$~kpc is higher than their stellar velocity dispersion measurements by $\sim 180$~\kms\ \citep[see Figure 4 in][]{ric11}. Both M87 and NGC~3311 remind us of the need for caution when using dynamical tracers to constrain the mass of a galaxy. We have carried out a dynamical analysis on the stars along the SE major axis at R~$\approx 40$ to 45~kpc from the center of M87. Even our most conservative interpretation finds the stellar velocity dispersion exceeding 400~\kms. The stellar kinematics are complex, with a strong asymmetry to positive velocity. Although speculative, the asymmetry may stem from a cooler, second component of stars superimposed on our field. All of this raises the question of what a rising velocity dispersion tells us about a galaxy. Is it a true reflection of the gravitational potential of the galaxy, the center of the galaxy cluster, or simply a snapshot of a dynamical system that has not reached equilibrium? The question of what our measurements are telling us about a given galaxy becomes more acute when the various dynamical tracers do not agree, as we have found in M87. Although groups have found good agreement between various tracers of mass in a wide range of galaxies \citep[e.g.,][]{coc09,mcn10}, we now know of many galaxies that show strong disagreement, and there is no a~priori reason that the GC, PNe and stellar kinematics must align at all radial positions. We know from simulations that the formation of BCGs appears very active \citep[e.g.,][]{del07a,rus09} and so this type of substructure and disagreement between different dynamical tracers should not be surprising and perhaps even expected. With a single field at this distance from the center of M87 we can not say much about the overall formation history of M87 beyond confirming that the center of the Virgo Cluster is still undergoing active assembly and the stellar halo exhibits a rising velocity dispersion profile. Further work on the observational front in necessary, yet the ability to constrain the dynamical state of the stars at these unprecedented radial distances points the way towards future exploration. In particular, a larger number of measurements of the kinematics of integrated starlight, taken at a broad range of locations, would be highly illuminating. | 14 | 3 | 1403.1581 |
1403 | 1403.1062_arXiv.txt | {}{We investigate the relation between 1D atmosphere models that rely on the mixing-length theory and models based on full 3D radiative hydrodynamic (RHD) calculations to describe convection in the envelopes of late-type stars.}{The adiabatic entropy value of the deep convection zone, $\ssbot$, and the entropy jump, $\Delta s$, determined from the 3D RHD models, were matched with the mixing-length parameter, $\amlt,$ from 1D hydrostatic atmosphere models with identical microphysics (opacities and equation-of-state). We also derived the mass mixing-length parameter, $\am$, and the vertical correlation length of the vertical velocity, $C\left[v_{z},v_{z}\right]$, directly from the 3D hydrodynamical simulations of stellar subsurface convection.}{The calibrated mixing-length parameter for the Sun is $\amlt^{\odot}\left(\ssbot\right)=1.98$. For different stellar parameters, $\amlt$ varies systematically in the range of $1.7-2.4$. In particular, $\amlt$ decreases towards higher effective temperature, lower surface gravity and higher metallicity. We find equivalent results for $\amlt^{\odot}\left(\Delta s\right)$. In addition, we find a tight correlation between the mixing-length parameter and the inverse entropy jump. We derive an analytical expression from the hydrodynamic mean-field equations that motivates the relation to the mass mixing-length parameter, $\am$, and find that it qualitatively shows a similar variation with stellar parameter (between $1.6$ and $2.4$) with the solar value of $\am^{\odot}=1.83$. The vertical correlation length scaled with the pressure scale height yields $1.71$ for the Sun, but only displays a small systematic variation with stellar parameters, the correlation length slightly increases with $\teff$.}{We derive mixing-length parameters for various stellar parameters that can be used to replace a constant value. Within any convective envelope, $\am$ and related quantities vary strongly. Our results will help to replace a constant $\amlt$.}{} | } In the past century, insights in various fields of physics led to a substantially more accurate interpretation and understanding of the processes taking place in the interior of celestial bodies. Astronomers can parameterize the conditions on the surface of stars with theoretical stellar atmosphere models, and with the theory of stellar structure and evolution, they are additionally capable to predict the complex development of stars. The radiated energy of cool stars, originating from the deeper interior because of nuclear burning in the center, is advected to the surface by convective motions in the envelope that are driven by negative buoyancy acceleration. At the thin photospheric transition region the large mean free path of photons allows them to escape into space, and the convective energy flux is abruptly released. To theoretically model this superadiabatic boundary domain of stars is challenging because of the nonlinear and nonlocal nature of turbulent subsurface convection and radiative transfer, and an analytical solution is a long-standing unresolved problem. To account for the convective energy transport, \citet{BohmVitense:1958p4822} formulated the mixing-length theory (MLT), which was initially proposed by \citet{Prandtl:1925} in analogy to the concept of the mean free path in the kinetic gas theory. In the framework of MLT, it is assumed that the heat flux is carried by convective elements for a typical distance before they dissolve instantaneously into the background. This distance is the so-called mixing-length, $\lmlt$, usually expressed in units of the pressure scale height, $\amlt=\lmlt/H_{P}$. The mixing-length parameter $\amlt$ is a priori unknown, hence it has to be calibrated, usually by matching the current radius and luminosity of the Sun by a standard solar model with a single depth-independent $\amlt^{\odot}$. This calibrated value for the Sun is then used for all stellar parameters. We recall that $\amlt^{\odot}$, in fact, corrects for all other shortcomings of the solar model, deficits in the equation-of-state (EOS), the opacities, or the solar composition. It therefore is no wonder that its numerical value \citep[typically around 1.7 to 1.9; e.g., see][]{Magic:2010p13816} varies with progress in these aspects and from code to code. In addition, MLT is a local and time-independent theory that effectively contains three additional, free parameters, and assumes symmetry in the up- and downflows, hence also in the vertical and horizontal direction. The actual formulation of MLT can also vary slightly \citep[e.g., see][]{Henyey:1965p15592,Mihalas:1970p21310,Ludwig:1999p7606}. Many attempts have been made to improve MLT, a substantial one being the derivation of a nonlocal mixing-length theory \citep{Gough:1977p8031,Unno:1985p22746,Deng:2006p22550,Grossman:1993p7953}. The standard MLT is a local theory, meaning that the convective energy flux is derived purely from local thermodynamical properties, ignoring thus any nonlocal properties (e.g., overshooting) of the flow. Nonlocal models are typically derived from the hydrodynamic equations, which are a set of nonlinear moment equations including higher order moments. To solve them, closure approximations are considered (e.g., diffusion approximation, anelatistic approximations, or introducing a diffusion length). Other aspects have also been studied: the asymmetry of the flow by a two-stream MLT model \citep{Nordlund:1976p13639}, the anisotropy of the eddies \citep{Canuto:1989p22737}, the time-dependence \citep{Xiong:1997p22562}, and the depth-dependence of $\amlt$ \citep{Schlattl:1997p1676}. While standard MLT accounts for only a single eddy size (which is $\lmlt$), \citet{Canuto:1991p6553} extended this to a larger spectrum of eddy sizes by including the nonlocal second-order moment \citep{Canuto:1996p4779}. The original Canuto-Mazzitelli theory -- also known as the\emph{ full spectrum turbulence} model -- used the distance to the convective region border as a proxy for the mixing-length; a later version \citep{Canuto:1992} re-introduced a free parameter resembling $\amlt$. These approaches are often complex, but so far, the standard MLT is still widely in use, and a breakthrough has not been achieved, despite all the attempts for improvements. In 1D atmosphere modeling, the current procedure is to assume a universal value of $1.5$ for the mixing-length parameter $\amlt$ \citep[see]{Gustafsson:2008p3814,Castelli:2004p4949}. For full stellar evolution models, the solar ``calibration'' yields values around $\sim1.7-1.9$ \citep[see, e.g.\ ]{Magic:2010p13816}. Since the value of the mixing-length parameter sets the convective efficiency and therefore changes the superadiabatic structure of stellar models, an accurate knowledge of $\amlt$ for different stellar parameters would be a first step in improving models in that respect. However, apart from the Sun, other calibrating objects are rare and data are much less accurate (see Sect.~\ref{sub:Comparison-with-observations} for an example), such as binary stars with well-determined stellar parameters. The mixing-length parameter can be deduced from multidimensional radiative hydrodynamic (RHD) simulations, where convection emerges from first principles \citep[e.g., see][]{Ludwig:1999p7606}. Over the past decades, the computational power has increased and the steady development of 3D RHD simulations of stellar atmospheres has established their undoubted reliability by manifold successful comparisons with observations \citep{Nordlund:1982p6697,Steffen:1989p18861,Ludwig:1994p18892,Freytag:1996p808,Stein:1998p3801,Nordlund:1990p6720,Nordlund:2009p4109}. The 3D RHD models have demonstrated that the basic picture of MLT is incorrect: there are no convective bubbles, but highly asymmetric convective motions. Nonetheless, an equivalent mixing-length parameter has been calibrated by \citet{Ludwig:1999p7606} based on 2D hydrodynamic models by matching the resulting adiabats with 1D MLT models \citep[see ][for the metal-poor cases]{Freytag:1999p7637}. The authors showed that $\amlt$ varies significantly with the stellar parameters (from $1.3$ to $1.8$), and also studied the impact of a variable $\amlt$ on a globular cluster \citep{Freytag:1999p7645}. In addition, \citet{Trampedach:2007p5614} applied a grid of 3D atmosphere models with solar metallicity to calibrate the mixing-length parameter (from 1.6 to 2.0), and the so-called mass mixing-length \citep{Trampedach:2011p5920}. In the present work we calibrate the mixing-length parameter with a 1D atmosphere code that consistently employs the identical EOS and opacity as used in the 3D RHD simulations (Sect.~\ref{sec:theoretical_models}). We present the resulting mixing-length parameter in Sect.~\ref{sec:Mixing-length}. We also determine the mass mixing-length -- the inverse of the logarithmic derivative of the unidirectional mass flux -- in Sect. \ref{sec:Mass-mixing-length}, and the vertical correlation length of the vertical velocity (Sect.~\ref{sec:velocity_correlation_length}) directly from the 3D atmosphere models. For the former quantity, we derive a relation from the hydrodynamic mean-field equations that demonstrates the relation to $\amlt$, which is further substantiated by our numerical results. Finally, we conclude in Sect.~\ref{sec:Conclusions}. | } We have calibrated the mixing-length parameter using realistic 3D RHD simulations of stellar surface convection by employing a 1D MLT stellar atmosphere code with identical microphysics. The calibration was achieved by varying the mixing-length parameter and matching the adiabatic entropy value of the deeper convection zone, $\ssbot$, or alternatively, matching the entropy jump, $\Delta s$. In both ways we found the mixing-length to decrease for higher $\teff$ and $\feh$, and lower $\logg$. The mixing-length varies in the range of $1.7-2.3$ for $\amlt\left(\ssbot\right)$ and $\sim1.8-2.4$ for $\amlt\left(\Delta s\right)$, and will lead to differences of up to $\pm20\,\%$ in $\amlt$ depending on the stellar mass. This changes the stellar interior structure by extending or shortening the depth of the convection zone and thus the stellar evolution; we intend to investigate in future studies how in detail a realistic $\amlt$ will impact basic stellar evolution predictions. Furthermore, we derived from the hydrodynamic mean field equations (for the first-time) a physically motivated connection of the mass mixing-length, which is the inverse of the vertical mass flux gradient, with the mixing-length. We determined the mass mixing-length parameter and found that it varies qualitatively similar to the mixing-length parameter in the range of $1.6-2.3$. The mass mixing-length parameter is also depth-dependent and decreases above the surface to lower values around $\sim0.5$, which agrees with previous findings from observations. Finally, the mass mixing-length parameter and mixing-length parameter strongly correlate with the logarithmic inverse of the entropy jump for different stellar parameters, that is $\amlt\sim-\ln\Delta s$. Finally, we also derived the vertical velocity correlation length, which features values similar to that of the mixing-length with approximately $\sim1.6-1.8$ of pressure scale height, but, the dependence with $\teff$ is inverted, meaning that the correlation length decreases with $\teff$. To summarize the importance of our work: we can finally remove the free parameters inherent in MLT and also avoid having to use solar calibrations for other stars. | 14 | 3 | 1403.1062 |
1403 | 1403.6018_arXiv.txt | The spectral study of the HESS J1745-290 high energy gamma-ray cut-off from the galactic center is compatible with a signal of Dark Matter (DM) annihilation or decay. If this is the case, a neutrino flux from that source is also expected. We analyze the neutrino flux predicted by DM particles able to originate the HESS J1745-290 gamma-rays observations. We focus on the electroweak and hadronic channels, which are favoured by present measurements. In particular, we study DM annihilating into $W^+W^-$ and $u\bar u$ with DM masses of $48.8$ and $27.9$ TeV respectively. We estimate the resolution angle and exposition time necessary to test the DM hypothesis as the origin of the commented gamma signal. | Different telescopes have observed Very High Energy (VHE) gamma-rays coming from the Galactic Center (GC), such as CANGAROO \cite{CANG}, VERITAS \cite{VER}, MAGIC \cite{MAG} or Fermi-LAT \cite{Vitale, ferm}. In this work, we will pay attention to the data collected by the HESS collaboration from the J1745-290 source during the years 2004, 2005, and 2006 \cite{Aha, HESS}. The variability of the IR and X-ray observations \cite{X} indicates a different emission mechanism for this part of the spectrum. In addition, one of the most characteristic features of the HESS J1745-290 data consists in a cut-off at several tens of TeVs. These spectral properties can be explained naturally by the photons produced by the annihilation or decay of Dark Matter (DM) particles. This interpretation was discussed from the very early days of the publication of the HESS data \cite{Bergstrom1,DMint} but it was concluded that the DM origin was disfavored \cite{DMint}. However, a recent study has shown that the observed data are well fitted as DM signal complemented by a diffuse background \cite{HESSfit}. Indeed, this background has a good motivation since VHE photons are also expected from radiative processes generated by particle acceleration in the neighborhood of the supermassive black hole Sgr A and the Sgr A East supernova. The analysis shows good agreement with DM annihilation or decay into $u\bar u$, $d\bar d$, $s\bar s$ and $t\bar t$ quark-antiquark channels and $W^+W^-$ and $ZZ$ boson channels. Leptonic and other quark-antiquark channels were excluded with $95.4\%$ confidence level. The background provided by the analysis is also compatible with the Fermi-LAT data from the IFGL J1745.6-2900 source observed during 25 months \cite{ferm}, which is spatially consistent with the HESS J1745-290 source \cite{Cohen}. In any case, the fundamental nature of this gamma-ray flux is still unclear. The entire VHE spectrum may be produced by particle propagation \cite{ferm,SgrA} in the vicinity of the commented supernova remnant and black hole, both located at the central region of our galaxy \cite{Atoyan,AN}. In addition, the emission region is quite compact since the signal is limited to a region of few tenths of degree \cite{HESS}. This feature is not consistent with dark halos simulated with non-baryonic cold DM, such as the standard NFW profile \cite{Navarro:1996gj}. It needs to be more compact as the ones produced when baryonic effects are taken into account. It has been argued that the baryonic gas falls to the inner part of the halo, modifying the gravitational potential and increasing the DM density in the center \cite{Blumenthal,Prada:2004pi}. This scenario is not completely accepted (read \cite{Romano} for example), but if it is correct, it has two important consequences. First, the sensitivity of indirect DM searches is reduced to a more compressed region; and second, the DM annihilating fluxes are enhanced by up to three orders of magnitude with respect to the standard NFW profile \cite{Prada:2004pi}. The HESS observations are in good agreement with these types of compressed dark halos. The DM particle that originate this spectrum needs to have a mass between $15 \; \text{TeV} \lesssim M \lesssim 110 \; \text{TeV}$ \cite{HESSfit}. This makes highly challenging to observe these particles in direct detection experiments or particle accelerators \cite{lab}. On the contrary, complementary cosmic rays analysis \cite{cosmics} from the GC and from other astrophysical objects are the most promising way to cross check the commented DM hypotheses. In particular, the analysis of neutrino fluxes from the same region can be determining. If DM annihilates or decays into Standard Model (SM) particles producing VHE gamma-rays photons, it has to produce also VHE neutrinos. Indeed, if the dark halo properties are adjusted to explain the HESS J1745-290 data, the neutrino flux is completely determined if one concrete annihilation or decay channel is assumed. This work is organised as follows: In Section II, we study the expected neutrino fluxes as indirect products of annihilating DM in the direction of the GC. Section III is devoted to discuss the flavor oscillation effects in this signal. In Section IV, we model the background of our analysis by taking into account the atmospheric neutrino flux observed by the IceCube experiment and we study the best configuration that may allow the detection of the corresponding neutrino signal associated with the HESS J$1745$-$290$ GC gamma-rays source. Finally, we summarize our main conclusions in Section V. \begin{figure}[bt] \begin{center} \epsfxsize=13cm \resizebox{8.8cm}{6.6cm} {\includegraphics{MCflux.pdf}} \caption {\footnotesize{ The gramma-ray ($\gamma$) and neutrino ($\nu_p$) fluxes from DM annihilating into $W^+W^-$ bosons, as generated by PYTHIA 8.135 and reported by \cite{Cirelli}.} } \label{MCfluxes} \end{center} \end{figure} | The operation of the IceCube neutrino telescope at the South Pole, together with several counterparts at the Nothern hemisphere, such as ANTARES and NT200 presently, or the future KM3NeT and GVD, are opening a new window in our knowledge of neutrino astronomy. \begin{figure}[t!] \begin{center} \epsfxsize=13cm \resizebox{8cm}{8cm} {\includegraphics{uAeff50t2ylinvsEmin.pdf}} \end{center} \begin{center} \epsfxsize=13cm \resizebox{8cm}{8cm} {\includegraphics{u06AefftexpvsEmin.pdf}} \caption {\footnotesize{ Confidence level contours associated to the observation of DM annihilating into the $u\bar u$ quark-antiquark channel at $1\sigma$ (dark), $2\sigma$, $3\sigma$, $4\sigma$, $5\sigma$ (white) confidence level. Top panel : The minimum energy cut is optimized around $1$ TeV depending on the resolution angle. The exposition time and effective area are fixed to the relation: $Af\equiv A_{\text{eff}}\times t_{\text{exp}}\simeq 100\, \text{m}^2\,\text{yr}$. Bottom panel: the angular field of view is fixed as $\theta=0.6^\circ$. In such a case, the possibility to detect the neutrino flux signal above the atmospheric background demands $Af\equiv A_{\text{eff}}\times t_{\text{exp}}\gtrsim 100\;\text{m}^2\text{yr}$.}} \label{usigma} \end{center} \end{figure} Indeed, the construction of KM3NeT will imply a new substantial improvement in sensitivity corresponding to a km$^3$ sized detector. On the other hand, radio and airshower detectors, such as ANITA and the Pierre Auger observatory are sensitive to neutrinos with even higher energies. The development of neutrino detectors have increased the interest for analysing the DM nature through the production of astrophysical neutrinos as its primary source. We have studied the prospective neutrino fluxes that should be originated by DM annihilating in the GC, in the case that the J1745-290 HESS high energy gamma-rays have this origin \cite{HESSfit}. The photon spectra is well fitted by different electroweak and hadronic channels. We have done a explicit analysis for $48.8$ TeV DM annihilating in $W^+W^-$ and $27.9$ TeV DM annihilating into $u\bar u$ channel. In these cases, the neutrino fluxes are completely determined by assuming that the DM region is localized as it is imposed by the gamma-rays analysis. We have estimated the best combinations of energy cuts, observation times and angular resolutions of a general high energy neutrino telescope. For this purpose, we have used IceCube atmospheric neutrino observations as background. In particular, the data collected with exposition time of $t_{\text{exp}}^{\nu_\mu}=359$ days and $t_{\text{exp}}^{\nu_e}=281$ days for the muon and electron neutrinos, respectively \cite{numu, nue}. We have found that for DM annihilating into the $W^+W^-$ boson channel, we need a resolution angle $0.18^\circ\lsim\theta\lsim0.72^\circ$ and low energy cut-off $818\,\text{GeV}\lsim E_{\text{min}}^\nu\lsim 1811$ GeV to get a signal between $5\sigma$ and $2\sigma$ with a minimum of 2 years of exposition time and a maximum of five years for a $50\;\text{m}^2$ of detector effective area. The mass associated with the $u\bar u$ annihilation channel is significantly smaller. It implies that the neutrino flux produced in this case is less energetic, and more difficult to discriminate from the background. It demands a higher angular resolution ($0.13^\circ\lsim\theta\lsim0.60^\circ$) and the energy cuts need to be smaller ($274\,\text{GeV}\lsim E_{\text{min}}^\nu\lsim 552$ GeV) in order to accumulate enough events. We have considered only track signal data by rejecting the muon background and taking into account the total number of events. For a binned analysis with a non-zero background and with a combined analysis of track and shower signatures, it could be possible to find better experimental configurations that should allow to detect neutrinos produced by heavy DM from the GC with worst resolution angle, smaller effective area or less exposition time. Recently, the IceCube collaboration have reported the observation of 28 high energy neutrinos over the range $30$ TeV $-\;1$ PeV at $4.1\sigma$ of confidence level, and $t_{\text{exp}}=662\text{ days}$ $(\simeq1.8\text{ years})$. Of these events, 5 are likely originated from the GC \cite{GCnu}. These neutrinos seem to have an astrophysical origin, but the spectrum and direction are not compatible with the signal studied in this work (the angular resolution in the muon track events is of $\theta\approx8^\circ$). The DM signal analyzed in this work may only account for a small part of the events, that will be more likely associated with an electroweak channel, as the $W^+W^-$ annihilating DM model. \vspace{0.5cm} {\bf Acknowledgements} We would like to thank Juande Zornoza and Carlos de los Heros for useful comments. This work has been supported by UCM predoctoral grant, MICINN (Spain) project numbers FIS 2008-01323, FIS2011-23000, FPA2011-27853-01 and Consolider-Ingenio MULTIDARK CSD2009-00064. | 14 | 3 | 1403.6018 |
1403 | 1403.3121_arXiv.txt | Axions are well motivated particles that could make up most or all of the dark matter if they have masses below 100 $\mu$eV. Microwave cavity techniques comprised of closed resonant structures immersed in solenoid magnets are sensitive to dark matter axions with masses of a few~$\mu$eV, but face difficulties scaling to higher masses. We present the a novel detector architecture consisting of an open, Fabry-P\'{e}rot resonator and a series of current-carrying wire planes, and demonstrate this technique with a search for dark matter axion-like particles called Orpheus. This search excludes dark matter axion-like particles with masses between 68.2 and 76.5 $\mu$eV and axion-photon couplings greater than $4\times10^{-7}~\mathrm{GeV}^{-1}$. We project that the fundamental sensitivity of this technique could be extended to be sensitive to couplings below $1\times10^{-15}~\mathrm{GeV}^{-1}$, consistent with the DFSZ model of QCD axions. | The axion is a pseudo-scalar particle predicted as a consequence to the Peccei-Quinn solution to the Strong CP problem~\cite{Peccei,Peccei_2,PhysRevLett.40.223,PhysRevLett.40.279}, and may comprise some or all of dark matter~\cite{Preskill1983127,Abbott1983,ipser-sikivie}. The axion has weak coupling to the electromagnetic interaction arising at loop order, whose Lagrange density may be written compactly as \begin{equation} {\cal L}_{a\gamma\gamma}=-g_{a\gamma\gamma}a\vec{E}\cdot\vec{B}~, \label{eqn:lagrangian} \end{equation} \noindent where $g_{a\gamma\gamma}$ is the axion-photon coupling strength, $a$ is the axion field, and $\vec{E}$,~$\vec{B}$ are the usual electric and magnetic fields. The expression in Eqn.~\ref{eqn:lagrangian} motivates the Axion Haloscope technique~\cite{PhysRevLett.51.1415} to detect dark matter axions. A typical Axion Haloscope consists of a closed microwave resonator immersed in a high static magnetic field, coupled to a low noise microwave receiver via the lowest frequency TM mode of the resonator. Dark matter axions passing through the magnetic field can convert into photons inside the cavity with enhanced probability when an electromagnetic resonance in the cavity is tuned to correspond to the frequency of the photons produced. Dark matter axions would be detected as excess power at this frequency, the expression for which can be derived from Eqn.~\ref{eqn:lagrangian} as \cite{PhysRevLett.80.2043} \begin{equation} P=\frac{2\pi\hbar^2 g^2_{a\gamma\gamma}\rho_{\rm DM}}{m^2_ac}\cdot f_{\gamma}\cdot \frac{1}{\mu_0}B^2V_{nlm}\cdot Q~. \label{eqn:axpow} \end{equation} \noindent Here the $m_a$, $f_{\gamma}$ denote the axion mass and frequency of the converted photon respectively and $\rho_{\rm DM}\approx 0.4~{\rm GeV/cc}$ is the local halo density of dark matter. The enhancement in the expected axion power due to its conversion in a resonant cavity is expressed in terms of the cavity quality factor $Q$. The effective volume of the cavity for coupling to a given resonant mode is \cite{Peng2000569} \begin{equation} V_{nlm}=\frac{\left(\int d^3\vec{x} \vec{E}(\vec{x})\cdot\vec{B}(\vec{x})\right)^2}{B^2\int d^3\vec{x} |\vec{E}|^2(\vec{x})}~, \label{eqn:Vnlm} \end{equation} \noindent where $\vec{B}(\vec{x})$ is the static magnetic field and $\vec{E}$ is the electric field of a normal resonant mode denoted by integers $n,~l,~m$. Numerous experiments based on this architecture have been constructed. Recently, the ADMX collaboration has demonstrated that microwave cavity experiments can be built with the sensitivity necessary to detect dark matter axions with masses in the range, 1.90--3.54 $\mu$eV~\cite{PhysRevLett.80.2043,PhysRevLett.104.041301} and coupling strength consistent with QCD predictions. Some models, however, predict the axion mass scale to be somewhat larger~\cite{PhysRevD.80.035024, Khlopov1999105,PhysRevD.85.105020}. Work is underway to extend experimental reach to larger axion masses, but the closed resonator detector design is difficult to extend to masses as large as 100~$\mu$eV~\cite{PhysRevD.64.092003}. Physically the size of a closed resonator must decrease in order to achieve higher resonant frequencies. This in turn decreases both the volume and $Q$ of the resonator, which both limits the sensitivity of experiments based on this architecture and presents a serious challenge to their scalability. We present a dark matter axion search technique which overcomes the fundamental limitations of closed resonator architectures at large axion masses by employing an open, Fabry-P\'{e}rot resonator as the detector volume. This technique is demonstrated by a prototype experiment named Orpheus. | We have presented a technique applicable to dark matter axion searches in the mass range 40--400 $\mu$eV, and demonstrated the technique with an experiment that searches for axion-like particles in the 68.2--76.5 $\mu$eV mass range, which is favored by some models of axion dark matter.\cite{PhysRevD.80.035024} Reasonable estimates based on the technology available, combined with the performance of a small scale implementation, suggest that experiments using this technique could be constructed to explore the majority of theoretically allowed axion-photon couplings. This work was supported in part by the U.S. Department of Energy under contract DE-SC0009800. | 14 | 3 | 1403.3121 |
1403 | 1403.3317_arXiv.txt | {We describe the drift field in thick depleted silicon sensors as a superposition of a one-dimensional {\it backdrop} field and various three-dimensional {\it perturbative} contributions that are physically motivated. We compute trajectories for the conversions along the field lines toward the channel and into volumes where conversions are confined by the perturbative fields. We validate this approach by comparing predictions against measured response distributions seen in five types of fixed pattern distortion features. We derive a quantitative connection between "tree ring" flat field distortions to astrometric and shape transfer errors with connections to measurable wavelength dependence -- as ancillary pixel data that may be used in pipeline analysis for catalog population. Such corrections may be tested on DECam data, where correlations between tree ring flat field distortions and astrometric errors -- together with their band dependence -- are already under study. Dynamic effects, including the brighter-fatter phenomenon for point sources and the flux dependence of flat field fixed pattern features are approached using perturbations similar in form to those giving rise to the fixed pattern features. These in turn provide drift coefficient predictions that can be validated in a straightforward manner. Once the three parameters of the model are constrained using available data, the model is readily used to provide predictions for arbitrary photo-distributions with internally consistent wavelength dependence provided for free.} | Using a physical model for a high resistivity, thick, fully depleted CCD sensor, we computed what we believe are intrinsic drift field line distortions and corresponding drift coefficients for points within the Si bulk that map to the saddle point loci that form the depth specific pixel boundaries according to $\delta \vec{x}_\perp(\vec{x}_0|\vec{x}_{sp})$. The drift coefficients are scaled appropriately to match laboratory characterization data that exhibit flat field response distortions, and consistent ancillary pixel data (up to six values per pixel) are computed. Ancillary pixel data -- when computed and utilized properly -- would smooth out flat field exposures, correct astrometric errors, and cancel pixel elongation and shape transfer effects (each to first order), and thereby remove dominant instrument signature features intrinsic to the raw data. Modifications to catalog population algorithms, image stacking, and multi-fit strategies are foreseen. We see the use of ancillary pixel data $\Delta I_{i,i+1,j,j+1}$, $\Delta P_{i,i+1,j,j+1}$ and $\Delta S_{i,i+1,j,j+1}$ as a scientifically prudent and affordable option: an alternative to masking off, or specially treating, the data from significant fractions of the focal plane.\footnote{For the current segmentation geometry, focal plane losses are roughly 10\%, if affected regions are all 20 pixels wide} Once best guess first order errors have been accounted for, residual flat field response distortions may be correlated against residual astrometric error fields to assess second order corrections of these types. | 14 | 3 | 1403.3317 |
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1403 | 1403.7534_arXiv.txt | The discovery of rapidly variable Very High Energy (VHE; $E >$ 100 GeV) $\gamma$-ray emission from 4C $+$21.35 (PKS\, 1222+216) by MAGIC on 2010 June 17, triggered by the high activity detected by the {\em Fermi} Large Area Telescope (LAT) in high energy (HE; $E >$ 100 MeV) $\gamma$-rays, poses intriguing questions on the location of the $\gamma$-ray emitting region in this flat spectrum radio quasar. We present multifrequency data of 4C $+$21.35 collected from centimeter to VHE during 2010 to investigate the properties of this source and discuss a possible emission model. The first hint of detection at VHE was observed by MAGIC on 2010 May 3, soon after a $\gamma$-ray flare detected by {\em Fermi}-LAT that peaked on April 29. The same emission mechanism may therefore be responsible for both the HE and VHE emission during the 2010 flaring episodes. Two optical peaks were detected on 2010 April 20 and June 30, close in time but not simultaneous with the two $\gamma$-ray peaks, while no clear connection was observed between the X-ray and $\gamma$-ray emission. An increasing flux density was observed in radio and mm bands from the beginning of 2009, in accordance with the increasing $\gamma$-ray activity observed by {\em Fermi}-LAT, and peaking on 2011 January 27 in the mm regime (230 GHz). We model the spectral energy distributions (SEDs) of 4C $+$21.35 for the two periods of the VHE detection and a quiescent state, using a one-zone model with the emission coming from a very compact region outside the broad line region. The three SEDs can be fit with a combination of synchrotron self-Compton and external Compton emission of seed photons from a dust torus, changing only the electron distribution parameters between the epochs. The fit of the optical/UV part of the spectrum for 2010 April 29 seems to favor an inner disk radius of $<$6 gravitational radii, as one would expect from a prograde-rotating Kerr black hole. | \label{intro} In the last few years flat spectrum radio quasars (FSRQs) have been established as a distinct Very High Energy (VHE) $\gamma$-ray blazar subclass. So far three FSRQs have been detected at $E$ $>$ 100 GeV: 3C 279 \citep{albert08}, 4C $+$21.35 \citep{MAGIC_discovery}, and PKS\, 1510$-$089 \citep{cortina12,abramowicz13}. These detections were surprising. The VHE emission from FSRQs is expected to be absorbed internally, if emitted within the broad line region (BLR), or externally, for sources located at redshifts where the emission is strongly attenuated by $\gamma$$\gamma$ pair production via interaction with the Extragalactic Background Light (EBL) photons. In addition, since FSRQs usually have their synchrotron peak at relatively low frequencies (i.e., infrared/optical bands rather than UV/X-ray), their corresponding inverse Compton peak should fall at photon energies less than 1 GeV, and thus a detection at VHE is not expected. The FSRQ 4C $+$21.35 (also known as PKS\, 1222$+$21) has a redshift of $z = 0.432$ \citep{osterbrock87} with a peculiar bent large-scale radio structure \citep{saikia93}. Very large apparent superluminal motion ($\beta_{\rm app} \sim$ 10--15) has been detected on milliarcsecond scales for sub-components of the jet \citep{jorstad01, homan01}. On the other hand, the ratio of the core-to-extended radio fluxes at GHz frequencies is of the order of unity; thus it is formally a ``lobe-dominated quasar'' \citep{kharb04, wang04}. In GeV $\gamma$-rays the source was in a quiescent state from the start of the {\em Fermi Gamma-ray Space Telescope} mission in 2008 August until 2009 September. After this period a gradually increasing flux was observed, up to an interval of flaring activity in the first half of 2010. In particular, 4C $+$21.35 underwent two very strong outbursts in 2010 April and June, observed by the Large Area Telescope (LAT) on board {\em Fermi} and composed of several major flares characterized by rise and decay timescales of the order of one day \citep{tanaka11}. During the second flaring activity, VHE emission from 4C $+$21.35 was detected with the MAGIC Cherenkov telescopes on 2010 June 17, with a flux doubling time of about 10 minutes \citep{MAGIC_discovery}. The simultaneous {\em Fermi}-LAT and Major Atmospheric Gamma Imaging Cherenkov (MAGIC) spectrum, corrected for EBL absorption, can be described by a single power-law with photon index $\Gamma_{\gamma}$ = 2.72 $\pm$ 0.34 between 3 GeV and 400 GeV, consistent with emission from a single component in the jet. The absence of a spectral cut-off for $E$ $<$ 130 GeV constrains the $\gamma$-ray emission region to lie outside the BLR, which would otherwise absorb the $\approx$ 10-20 GeV photons by $\gamma\gamma\rightarrow e^\pm$ production when these $\gamma$-rays pass through the intense circum-nuclear photon fields provided by the BLR itself. At the same time, the rapid VHE variability observed suggests an extremely compact emission region, with size $R \le c t_{\rm var} \delta_D/(1+z) \sim 10^{15}\ (\delta_D/80)\ (t_{\rm var}/10\ {\rm minutes})$\ cm where $t_{\rm var}$ is the variability timescale and $\delta_D$ is the Doppler factor. If the blob takes up the entire cross section of the jet, it implies that the emitting region is at a distance $r \sim R/\theta_{\rm open} \sim 5.7\times10^{16} (\delta_D/80)\ (t_{\rm var}/10\ {\rm minutes})$\ ($\theta_{\rm open}/1\deg)^{-1}$\ cm, where $\theta_{\rm open}$ is the half-opening angle of the jet. Even for a highly relativistic jet with $\delta_D \sim 100$, the location of the emission region should be well within the BLR radius for 4C +21.35, likely $R_{\rm BLR}\approx 2\times10^{17}$\ cm \citep{tanaka11}. Different models have been proposed to explain the unusual behavior of 4C $+$21.35. A very narrow jet can preserve variability at the pc scale, but the likelihood of being in the beam of such a thin jet is small, unless there were many narrow jets, as in a jets-within-jet/mini-jets scenario \citep{gt08,gub09,tav11}. An alternative model is a compact emission region at the pc scale responsible for the emission at higher energies, with a second zone either inside or outside the BLR to complete the modeling of the emission at lower energies \citep{tav11}. The compact emission sites at the pc scale could be due to self-collimating jet structures \citep{nal12}, where the magnetic field dominates the energy density, or to turbulent cells \citep[e.g.,][]{nal11,mj10}. Another possibility is that the acceleration of ultra-high energy cosmic rays protons in the inner jet leads to an outflowing beam of neutrons that deposit their energy into ultra-relativistic pairs that radiate VHE synchrotron emission at the pc scale \citep{dmt12}, with associated neutrino production. Even more exotic scenarios have been proposed, such as photons produced inside the BLR, tunneling through it via photon to axion-like particle oscillations \citep{tavecchio12}. In this paper, we present the multifrequency data of 4C $+$21.35 collected from radio to VHE during 2010, and discuss a possible emission model for this source. A summary of the complete multiwavelength data of 4C $+$21.35 presented in this paper and the relative facilities can be found in Table \ref{facilities}. The paper is organized as follows: in Sections 2 and 3 we briefly report the LAT and MAGIC data analysis and results, respectively. In Section 4 we report the result of {\em Swift} optical/UV/X-ray observations. Optical data collected by the Abastumani, ATOM, Catalina, Crimean, KVA, Steward, and St. Petersburg observatories are presented in Section 5. In Section 6 we present the radio and mm data collected by the Medicina, UMRAO, MOJAVE, OVRO, F-GAMMA, Mets\"ahovi, and SMA facilities. In Section 7 we discuss the light curves behavior and the spectral energy distribution (SED) modeling of three different epochs, and finally we draw our conclusions in Section 8. Throughout the paper, a $\Lambda$ CDM cosmology with $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda} = 0.73$, and $\Omega_{m} = 0.27$ is adopted. The corresponding luminosity distance at $z = 0.432$ is $d_L = 2370$\ Mpc, and 1 arcsec corresponds to a projected size of 5.6 kpc. \begin{table}[!hhh] \caption{Observatories Contributing to the Presented Data Set of 4C $+$21.35 at Different Frequencies.} \begin{center} \begin{tabular}{ccc} \hline \hline \label{facilities} Waveband & Observatory & Frequency/Band \\ \hline Radio & SMA & 230 GHz \\ & Mets\"ahovi & 37 GHz \\ & VLBA (MOJAVE) & 15 GHz \\ & OVRO & 15 GHz \\ & UMRAO & 8.0, 14.5 GHz \\ & Medicina & 5, 8 GHz \\ & F-GAMMA & 2.6, 4.8, 8.4, 10.5, 14.6, 23.1, 32, 86.2, 142.3 GHz \\ \hline Optical & Abastumani & $R$ \\ & ATOM & $R$ \\ & Catalina & $V$ \\ & Crimean & $R$ \\ & KVA & $R$ \\ & St. Petersburg & $R$ \\ & Steward & $V$ \\ & {\em Swift}-UVOT & $v$, $b$, $u$ \\ \hline UV & {\em Swift}-UVOT & $w1$, $m2$, $w2$ \\ \hline X-rays & {\em Swift}-XRT & 0.3--10 keV \\ & {\em Swift}-BAT & 15--50 keV \\ \hline HE $\gamma$-rays & {\em Fermi}-LAT & 0.1--300 GeV \\ VHE $\gamma$-rays & MAGIC & 70 GeV--5 TeV \\ \hline \hline \end{tabular} \end{center} \end{table} | \begin{enumerate} \item The $\gamma$-ray flares in 2010 April and June cannot have originated from inside the BLR, at least not without invoking some unusual particle transport mechanism \citep{dmt12,tavecchio12}. \item There is some evidence for a rapidly-spinning prograde BH based on the optical emission. \item The two flaring states and the quiescent state can be modeled by varying only the electron distribution for the source. \end{enumerate} The last result, modeling the source by varying only the electron distribution, has also been found for the blazar PKS\,0537$-$441 \citep{dammando13_0537}. This conclusion is much stronger for PKS\,0537$-$441, since the optical continuum of PKS\,0537$-$441 is not disk-dominated, making its modeling more constraining. Nonetheless, there are clearly sources for which a change in the electron distribution is not sufficient to explain the difference between flaring and quiescent states. For example, to model a strong optical-near infrared flare from PKS\,0208$-$512 with no counterpart in $\gamma$-rays required changing the magnetic field strength \citep{chatterjee13}. Rotation in polarization angles coincident with flares has been observed before in the blazars BL Lac \citep{marscher08}, PKS\,1510$-$089 \citep{marscher10,orienti13}, and 3C 279 \citep{abdo10_3c279}. They could be caused by a sudden realignment in the magnetic field due to shock compression, or a curved trajectory taken by the flaring region. A slight increase of the degree of optical polarization but no significant rotation of the polarization angle was observed at the time of the 2010 June HE and VHE flare. The object 4C~+21.35 continues to challenge our understanding of blazar emission mechanisms and the location of the emitting region. Multi-wavelength observations have complemented previous LAT and MAGIC observations to give a more complete picture for this source, although many outstanding questions remain. | 14 | 3 | 1403.7534 |
1403 | 1403.5531_arXiv.txt | \pagenumbering{arabic} The recent Planck results \cite{ade} have indicated that the cosmological perturbations in the Cosmic Microwave Background (CMB) radiation are nearly gaussian and of the adiabatic type. If one insists in assuming that these scalar perturbations are to be ascribed to single-field model of inflation \cite{lr}, the data put severe constraints restriction on the inflationary parameters. In particular, the Planck results have strengthened the upper limits on the tensor-to-scalar ratio, $r< 0.12 $ at 95\% C.L., disfavouring many inflationary models. In particular, the simplest quadratic chaotic model has been excluded at about 95\% C.L. Among the inflationary models discussed by the Planck collaboration is the Starobinsky $(R + R^2)$ theory, first presented in Refs. \cite{star} (see also Ref. \cite{Whitt:1984pd}). The Starobinsky model is a model that leads to a quasi-de Sitter phase and it is described by the Lagrangian \be S_{\rm S}=\frac{M_{\rm p}^2}{2}\int {\rm d}^4 x\sqrt{-g} \ \left(R+\frac{1}{6M^2} R^2\right), \label{star} \ee where $M_{\rm p}$ is the reduced Planck mass. This theory does not describe only the GR degrees of freedom, {\it i.e.} the helicity-2 massless graviton, but in addition it propagates a scalar degree of freedom usually called "scalaron". The later is hidden in the action (\ref{star}) and can be revealed in the so-called linear representation, where one writes the action in the equivalent form \cite{Whitt:1984pd} \be \label{R2} S_{\rm S}=\int {\rm d}^4 x\sqrt{-g} \ \left(\frac{M_{\rm p}^2}{2}R + \frac{M_p}{M} R\psi- 3 \psi^2\right). \ee After integrating out the field $\psi$, one gets back the original theory (\ref{star}). However, the action (\ref{R2}) is written in a Jordan frame and it can be expressed in Einstein frame after the conformal transformation \be g_{\mu\nu}\to e^{-\sqrt{2/3} \phi/M_{\rm p}}g_{\mu\nu}=\left(1+\frac{2\psi}{M M_{\rm p}}\right)^{-1} g_{\mu\nu} \ee is performed. Then, we get the equivalent scalar field version of the Starobinsky model \be S_{\rm S}=\int {\rm d}^4 x\sqrt{-g}\left[\frac{M_{\rm p}^2}{2}R-\frac{1}{2}\partial_\mu\phi\partial^\mu \phi-\frac{3}{4} M_{\rm p}^2 M^2\left(1-e^{-\sqrt{\frac{2}{3}}\phi/M_{\rm p}}\right)^2\right]. \label{R3} \ee There is a plateau in the scalar potential for large values of $\phi$ where slow-roll inflation can be realised with a quasi-de Sitter phase driven by a vacuum energy \be V_{\rm S}=\frac{3}{4}M_{\rm p}^2 M^2. \label{VS} \ee The normalization of the CMB anisotropies fixes $M\approx 10^{-5}M_{\rm p}$. In addition, the scalar tilt $n_S$ and tensor-to-scalar ratio $r$ turns out to be \be \label{pred} n_S-1 \approx -\frac{2}{N},\, \,\,r\approx \frac{12}{N^2}. \ee Note that $r$ has an addition $1/N$ suppression with respect to $n_S$. Although this model looks quite ad hoc at the theoretical level, it is perfect agreement with the Planck data, basically due to an additional $1/N$ suppression ($N$ being the number of e-folds till the end of inflation) of $r$ with respect to the prediction for the scalar spectral index $n_S$. For this reason, there has been a a renewed interest on the Starobinsky model, with particular emphasis on its supergravity extensions\cite{KL,sugrastaro,ENO,FKLP,FKR,FKD,KT}, along the lines originated in Refs. \cite{fv,Cecotti:1987sa,CFPS}. This positive attitude versus the Starobinsky model has dramatically changed with the recent release of the measurement of the tensor modes from large angle CMB B-mode polarization by BICEP2 \cite{bicep}, implying a tensor-to-scalar ratio \be r=0.2^{+0.07}_{-0.05}. \ee Putting aside the tension with the Planck data, this result (if confirmed) puts inflation on a ground which is firmer than ever. On the other side, it is in contradiction with the predictions (\ref{pred}) of the Starobinsky model. The goal of this paper is to show that this is not necessarily true: the contradiction with the tensor modes data disappear if one embeds the Starobinsky model in supergravity and identifies the inflaton field with the imaginary part of the chiral multiplet in the dual formulation of the model (instead of the the real part of it, as done in all the literature so far). We dub this version of the Starobinsky theory the ``Imaginary Starobinsky model" and show that it basically resembles the quadratic chaotic model during inflation \cite{chaotic} (for recent reviews, see \cite{Kallosh:2014ona}) once the coupling to matter is considered (a necessary condition to allow reheating in the model). It is nice that just embedding the Starobinky model into supergravity can make it in agreement with the data. Recently in \cite{KLr} it was shown that the simplest proposal in the standard supersymmetric Starobinsky model to identify the axion $b$, partner of the scalaron, with inflaton does not work and a drastic modification to the theory must be made if the $b$ field is responsible for the inflation. We show that a plausible modification can be made which naturally leads to $b$ inflation. The paper is organized as follows. In section 2 we recall the basics of the embedding of the Starobinsky model in supergravity, as done in the literature so far. In section 3 we describe our proposal to identify the inflaton with the partner of the ``scalaron" rather than the scalaron itself. Section 4 contains the main points about the imaginary Starobinsky model. Section 5 contains our conclusions. | In this paper we have reconsidered the prediction of the supersymmetric Starobinky model of inflation. In its dual formulation, although the real part of the chiral multiplet cannot generate enough tensor modes as an inflaton, its imaginary part does if appropriate couplings to matter are introduced. While its non-supersymmetric version seems to be ruled out by the recent BICEP2 data on the amount of tensor modes, we have shown that the field space of the supersymmetric theory contains inflationary directions which are in agreement with the current data once appropriate couplings to matter are considered. The reason is that, along this imaginary direction and once the couplings to matter are considered, the model may become the chaotic single-field model with a quadratic potential. | 14 | 3 | 1403.5531 |
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1403 | 1403.5688_arXiv.txt | Air-shower measurements in the {primary} energy range beyond 10\,TeV can be used to address {important} questions of astroparticle and particle physics. The most prominent among these questions are the search for the origin of charged Galactic cosmic rays and {the so-far little understood transition from Galactic to extra-galactic cosmic rays. A very promising avenue towards answering these fundamental questions is the construction of an air-shower detector with sufficient sensitivity for gamma-rays to {identify} the accelerators and large exposure to achieve accurate spectroscopy of local cosmic rays.} With the new ground-based large-area (up to 100\,km$^2$) wide-angle ({$\Omega\,\sim$0.6--0.85\,sr}) air-shower detector concept HiSCORE (Hundred*i Square-km Cosmic ORigin Explorer), we aim at exploring the cosmic ray and gamma-ray sky (accelerator-sky) in the energy range from {few 10s of TeV} to 1\,EeV {using the non-imaging air-Cherenkov detection {technique}}. The full detector simulation is presented here. The resulting sensitivity of a HiSCORE-type detector to gamma-rays will extend the energy range so far accessed by other experiments beyond energies of 50\,--100\,TeV, thereby opening up the ultra high energy gamma-ray (UHE gamma-rays, E$>$10\,TeV) observation window. | \label{hiscore_introduction} The current knowledge about the origin of cosmic rays has been accumulated following two different approaches: (i) by measuring in detail the energy spectrum and mass composition of the local cosmic-ray population and (ii) by gamma-ray ($E>100$~MeV) observations of both individual astrophysical objects as well as the diffuse emission from the interstellar medium. Both approaches provide complementary information/constraints on the most relevant quantities: e.g. the measurement of spallation products and cosmo-genic nuclei provides information on the energy dependence of cosmic-ray transport and the escape time of cosmic rays out of the Galaxy. Gamma-ray observations constrain the spatial distribution and properties of the cosmic ray accelerators and the density of cosmic rays in the interstellar medium. Cosmic-ray measurements through air-shower techniques are the only means to collect sufficient event statistics to measure cosmic rays at energies close to the knee ($\approx 3\times 10^{15}~$eV) in the all-particle energy spectrum. The traditional air-shower detectors sample the lateral density function (LDF) of secondary particles {or photons} on the ground. Given the large intrinsic fluctuations in the shower development and {that only a small fraction of the particles are sampled} ($\approx 10^{-4}$), the energy resolution and sensitivity to different primary particles is rather limited. Combining detection of different components of the air shower as e.g. realised in the KASCADE air shower field \cite{2004ApJ...608..865A}, improves the situation considerably but suffers from limited collection area. Established techniques to follow the longitudinal air shower development include muon tracking, air Cherenkov, and air fluorescence observations. The latter technique has been realised quite early and remains one of the most sensitive techniques at ultra-high energies \cite[Linsley, Fly's eye, HiRes, Pierre-Auger Observatory, Telescope-Array, see][and references therein]{2011ARA&A..49..119K}. The non-imaging air Cherenkov technique measures the arrival time and the LDF of the Cherenkov photons in the air shower front. This technique is sensitive to the longitudinal air shower development (mainly position of the shower maximum) as demonstrated with e.g. Themistocle \cite{1993APh.....1..341T}, AIROBICC \cite{1995APh.....3..321K}, Blanca \cite{2001APh....15...49F}, Tunka \cite{2012NIMPA.692...98B}, Jakutsk \cite{1986NIMPA.248..224D}. The longitudinal air shower development is sensitive to the initial particle species. Both techniques allow a comparably good energy resolution which suffers less from the fluctuations and the limited sampling. A number of new approaches for air shower detection have been proposed and partially tested including long-wavelength {(MHz)} radio measurements\footnote{the dominant emission processes are geo-synchrotron and charge separation} \cite[see e.g.][]{2012NIMPA.662S..72H}, and molecular Bremsstrahlung emission at {GHz frequencies} \cite{2008PhRvD..78c2007G,2013EPJWC..5308010S}. In the sense of shower-front sampling, the long-wavelength radio observations are comparable to the air Cherenkov technique while the molecular Bremsstrahlung has {analogies} to the air fluorescence (mostly isotropic emission) and {would allow} for imaging of the air shower development. For approach (i) -- spectroscopy and measurement of chemical composition of cosmic rays in the energy range from below the knee to the ankle ($10^{18}$~eV) -- the air Cherenkov approach appears to be {among} the best choice{{s}}, {considering its good energy resolution of the order of 10\,\%, and a typical resolution of the shower maximum of the order of 30\,g/cm$^2$ \cite{2012NIMPA.692...98B}. See also Section~\ref{section:air_shower_reconstruction}}. For approach (ii) -- gamma-ray observations -- the currently most succesfull technique is the imaging air Cherenkov technique with multiple telescopes (imaging air Cherenkov telescopes: IACTs). A large array of IACTs is currently under design to achieve a ten-fold improvement in flux sensitivity {as compared to current generation instruments}: the Cherenkov Telescope Array \cite[CTA, see][]{2011ExA....32..193A}. Nevertheless CTA is designed to achieve optimum sensitivity at TeV energies and will suffer from its limited collection area at energies {beyond 100\,TeV.} The non-imaging air Cherenkov technique allows to extend the collection area to {several} square kilometers with a moderate number of read-out channels\footnote{there are complementary approaches using large field of view cameras which would allow to increase the spacing of individual telescopes \cite{2008NIMPA.588...48R}}. In combination with the demonstrated good angular resolution in non-imaging Cherenkov air shower arrays, a multi-km$^2$ array with good sensitivity above 10 TeV appears feasible and is explored here. { With the Hundred*i Square-km Cosmic ORigin Explorer HiSCORE, we want to cover both approaches (i) and (ii) described above. A central question will be the search for the elusive pevatrons \cite{2007ApJ...665L.131G}, the accelerators of cosmic rays up to the PeV energy regime. } For more details on physics topics for HiSCORE, see \cite[][]{2011AdSpR..48.1935T} and references therein. | \label{hiscore_summary} We have described the HiSCORE concept for a large-area wide-angle air shower experiment, {based on an array of non imaging light collecting detector stations}. A comprehensive simulation of the detector was performed including all relevant components (atmosphere, light-collection with Winston cones, photomultiplier, pulse shaping with afterpulses, station trigger). The resulting effective areas for various primary particles and the expected trigger rates for background have been calculated for different assumptions on the trigger threshold. The complete Monte Carlo simulation of the HiSCORE detector concept shows that such a non-imaging air-Cherenkov detector will be sufficiently sensitive to survey a large fraction ($\pi~\mathrm{sr}$) of the sky for gamma-ray sources above {10--50\,TeV (depending on the final array layout)} at an energy flux level of a few $10^{-13}$~ergs/cm$^2$s. This sensitivity is comparable to the planned next-generation Cherenkov telescope array (CTA) at lower energies, effectively extending the sensitive energy range into the UHE gamma-ray regime. Furthermore, HiSCORE will provide high-statistics measurements of cosmic ray spectra and composition above 100\,TeV primary energy, covering the energy range of transition between Galactic and extragalactic origin of cosmic rays, and up to 10$^{18}$\,eV. An engineering array with $1-2~\mathrm{km}^2$ is planned for deployment {2014/2015}, aiming at proof-of-principle measurements and first physics results. | 14 | 3 | 1403.5688 |
1403 | 1403.0587_arXiv.txt | In an equilibrium axisymmetric galactic disc, the mean galactocentric radial and vertical velocities are expected to be zero everywhere. In recent years, various large spectroscopic surveys have however shown that stars of the Milky Way disc exhibit non-zero mean velocities outside of the Galactic plane in both the Galactocentric radial and vertical velocity components. While radial velocity structures are commonly assumed to be associated with non-axisymmetric components of the potential such as spiral arms or bars, non-zero vertical velocity structures are usually attributed to excitations by external sources such as a passing satellite galaxy or a small dark matter substructure crossing the Galactic disc. Here, we use a three-dimensional test-particle simulation to show that the global stellar response to a spiral perturbation induces both a radial velocity flow and non-zero vertical motions. The resulting structure of the mean velocity field is qualitatively similar to what is observed across the Milky Way disc. We show that such a pattern also naturally emerges from an analytic toy model based on linearized Euler equations. We conclude that an external perturbation of the disc might not be a requirement to explain all of the observed structures in the vertical velocity of stars across the Galactic disc. Non-axisymmetric internal perturbations can also be the source of the observed mean velocity patterns. | The Milky Way has long been known to possess spiral structure, but studying the nature and the dynamical effects of this structure has proven to be elusive for decades. Even though its fundamental nature is still under debate today, it has nevertheless started to be recently considered as a key player in galactic dynamics and evolution (e.g., Antoja et al.~2009; Quillen et al.~2011; L\'epine et al.~2011; Minchev et al.~2012; Roskar et al.~2012 for recent works, or Sellwood~2013 for a review). However, zeroth order dynamical models of the Galaxy still mostly rely on the assumptions of a smooth time-independent and axisymmetric gravitational potential. For instance, recent determinations of the circular velocity at the Sun's position and of the peculiar motion of the Sun itself all rely on the assumption of axisymmetry and on minimizing the non-axisymmetric residuals in the velocity field (Reid et al.~2009; McMillan \& Binney 2010; Bovy et al.~2012; Sch\"onrich~2012). Such zeroth order assumptions are handy since they allow us to develop dynamical models based on a phase-space distribution function depending only on three isolating integrals of motion, such as the action integrals (e.g., Binney~2013; Bovy \& Rix~2013). Actually, an action-based approach does not necessarily have to rely on the axisymmetric assumption, as it is also possible to take into account the main non-axisymmetric component (e.g., the bar, see Kaasalainen \& Binney 1994) by modelling the system in its rotating frame (e.g., Kaasalainen 1995). However the other non-axisymmetric components such as spiral arms rotating with a different pattern speed should then nevertheless be treated through perturbations (e.g., Kaasalainen~1994; McMillan~2013). The main problem with such current determinations of Galactic parameters, through zeroth order axisymmetric models, is that it is not clear that assuming axisymmetry and dynamical equilibrium to fit a benchmark model does not bias the results, by e.g. forcing this benchmark model to fit non-axisymmetric features in the observations that are not present in the axisymmetric model itself. This means that the residuals from the fitted model are not necessarily representative of the true amplitude of non-axisymmetric motions. In this respect, it is thus extremely useful to explore the full range of possible effects of non-axisymmetric features such as spiral arms in both fully controlled test-particle simulations as well as self-consistent simulations, and to compare these with observations. With the advent of spectroscopic and astrometric surveys, observational phase-space information for stars in an increasingly large volume around the Sun have allowed us to see more and more of these dynamical effect of non-axisymmetric components emerge in the data. Until recently, the most striking features were found in the solar neighbourhood in the form of moving groups, i.e. local velocity-space substructures shown to be made of stars of very different ages and chemical compositions (e.g., Chereul et al. 1998, 1999; Dehnen~1998; Famaey et al.~2005, 2007, 2008; Pomp\'eia et al.~2011). Various non-axisymmetric models have been argued to be able to represent these velocity structures equally well, using transient (e.g., De Simone et al. 2004) or quasi-static spirals (e.g., Quillen \& Minchev 2005; Antoja et al. 2011), with or without the help of the outer Lindblad resonance from the central bar (e.g., Dehnen 2000; Antoja et al. 2009; Minchev et al. 2010; McMillan 2013; Monari et al. 2013). The effects of non-axisymmetric components have also been analyzed a bit less locally by Taylor expanding to first order the planar velocity field in the cartesian frame of the Local Standard of Rest, i.e. measuring the Oort constants $A$, $B$, $C$ and $K$ (Kuijken \& Tremaine~1994; Olling \& Dehnen~2003), a procedure valid up to distances of less than 2~kpc. While old data were compatible with the axisymmetric values $C=K=0$ (Kuijken \& Tremaine~1994), a more recent analysis of ACT/Tycho2 proper motions of red giants yielded $C = -10 \, {\rm km}\,{\rm s}^{-1}\,{\rm kpc}^{-1}$ (Olling \& Dehnen~2003). Using line-of-sight velocities of 213713 stars from the RAVE survey (Steinmetz et al. 2006; Zwitter et al. 2008; Siebert et al. 2011a; Kordopatis et al. 2013), with distances $d<2 \,$kpc in the longitude interval $-140^\circ < l < 10^\circ$, Siebert et al. (2011b) confirmed this value of $C$, and estimated a value of $K= +6\,{\rm km}\,{\rm s}^{-1}\,{\rm kpc}^{-1}$, implying a Galactocentric radial velocity\footnote{In this paper, {\it 'radial velocity'} refers to the Galactocentric radial velocity, not to be confused with the line-of-sight (l.o.s.) velocity.} gradient of $C+K = \partial V_R / \partial R \simeq - 4\,{\rm km}\,{\rm s}^{-1}\,{\rm kpc}^{-1}$ in the solar suburb (extended solar neighbourhood, see also Williams et al. 2013). The projection onto the plane of the mean line-of-sight velocity as a function of distance towards the Galactic centre ($|l|<5^\circ$) was also examined by Siebert et al. (2011b) both for the full RAVE sample and for red clump candidates (with an independent method of distance estimation), and clearly confirmed that the RAVE data are not compatible with a purely axisymmetric rotating disc. This result is not owing to systematic distance errors as considered in Binney et al. (2013), because the {\it geometry} of the radial velocity flow cannot be reproduced by systematic distance errors alone (Siebert et al.~2011b; Binney et al.~2013). Assuming, to first order, that the observed radial velocity map in the solar suburb is representative of what would happen in a razor-thin disc, and that the spiral arms are long-lived, Siebert et al. (2012) applied the classical density wave description of spiral arms (Lin \& Shu 1964; Binney \& Tremaine 2008) to constrain their parameters in the Milky Way. They found that the best-fit was obtained for a two-armed perturbation with an amplitude corresponding to $\sim 15$\% of the background density and a pattern speed $\Omega_P \simeq 19 \,{\rm Gyr}^{-1}$, with the Sun close to the 4:1 inner ultra-harmonic resonance (IUHR). This result is in agreement with studies based on the location of moving groups in local velocity space (Quillen \& Minchev~2005; Antoja et al.~2011; Pomp\'eia et al.~2011). This study was advocated to be a useful first order benchmark model to then study the effect of spirals in three dimensions. In three dimensions, observations of the solar suburb from recent spectroscopic surveys actually look even more complicated. Using the same red clump giants from RAVE, it was shown that the mean {\it vertical} velocity was also non-zero and showed clear structure suggestive of a wave-like behaviour (Williams et al.~2013). Measurements of line-of-sight velocities for 11000 stars with SEGUE also revealed that the mean vertical motion of stars reaches up to 10~km/s at heights of 1.5~kpc (Widrow et al.~2012), echoing previous similar results by Smith et al. (2012). This is accompanied by a significant wave-like North-South asymmetry in SDSS (Widrow et al.~2012; Yanny \& Gardner~2013). Observations from LAMOST in the outer Galactic disc (within 2~kpc outside the Solar radius and 2~kpc above and below the Galactic plane) also recently revealed (Carlin et al.~2013) that stars above the plane exhibit a net outward motion with downward mean vertical velocities, whilst stars below the plane exhibit the opposite behaviour in terms of vertical velocities (moving upwards, i.e. towards the plane too), but not so much in terms of radial velocities, although slight differences are also noted. There is thus a growing body of evidence that Milky Way disc stars exhibit velocity structures across the Galactic plane in {\it both} the Galactocentric radial and vertical components. While a global radial velocity gradient such as that found in Siebert et al. (2011b) can naturally be explained with non-axisymmetric components of the potential such as spiral arms, such an explanation is {\it a priori} less self-evident for vertical velocity structures. For instance, it was recently shown that the central bar cannot produce such vertical features in the solar suburb (Monari et al. 2014). For this reason, such non-zero vertical motions are generally attributed to vertical excitations of the disc by external means such as a passing satellite galaxy (Widrow et al.~2012). The Sagittarius dwarf has been pinpointed as a likely culprit for creating these vertical density waves as it plunged through the Galactic disc (Gomez et al.~2013), while other authors have argued that these could be due to interaction of the disc with small starless dark matter subhalos (Feldmann \& Spolyar~2013). Here, we rather investigate whether such vertical velocity structures can be expected as the response to disc non-axisymmetries, especially spiral arms, in the absence of external perturbations. As a first step in this direction, we propose to qualitatively investigate the response of a typical old thin disc stellar population to a spiral perturbation in controlled test particle orbit integrations. Such test-particle simulations have revealed useful in 2D to understand the effects of non-axisymmetries and their resonances on the disc stellar velocity field, including moving groups (e.g., Antoja et al.~2009, 2011; Pomp\'eia et al.~2011), Oort constants (e.g., Minchev et al. 2007), radial migrations (e.g., Minchev \& Famaey~2010), or the dip of stellar density around corotation (e.g., Barros et al.~2013). Recent test-particle simulations in 3D have rather concentrated on the effects of the central bar (Monari et al.~2013, 2014), while we concentrate here on the effect of spiral arms, with special attention to mean vertical motions. In Sect. 2, we give details on the model potential, the initial conditions and the simulation technique, while results are presented in Sect. 3, and discussed in comparison with solutions of linearized Euler equations. Conclusions are drawn in Sect.~4. | In recent years, various large spectroscopic surveys have shown that stars of the Milky Way disc exhibit non-zero mean velocities outside of the Galactic plane in both the Galactocentric radial component and vertical component of the mean velocity field (e.g., Siebert et al. 2011b; Williams et al. 2013; Carlin et al. 2013). While it is clear that such a behaviour could be due to a large combination of factors, we investigated here whether spiral arms are able to play a role in these observed patterns. For this purpose, we investigated the orbital response of a test population of stars representative of the old thin disc to a stable spiral perturbation. This is done using a test-particle simulation with a background potential representative of the Milky Way. We found non-zero velocities both in the Galactocentric radial and vertical velocity components. Within the rotating frame of the spiral pattern, the location of these non-zero mean velocities in both components are stable over time, meaning that the response to the spiral perturbation is stable. Within corotation, the mean $\langle v_R \rangle$ is negative within the arms (mean radial motion towards the Galactic centre) and positive (radial motion towards the anticentre) between the arms. Outside corotation, the pattern is reversed, as expected from the Lin-Shu density wave theory (Lin \& Shu 1964). On the other hand, even though the spiral perturbation of the potential is very thin, the radial velocity flow is still strongly affected above the Galactic plane. Up to five times the scale-height of the spiral potential, there are no strong asymmetries in terms of radial velocity, but above these heights, the trend in the radial velocity flow is reversed. This means that asymmetries could be observed in surveys covering different volumes above and below the Galactic plane. Also, forthcoming surveys like Gaia, 4MOST, WEAVE will be able to map this region of the disc of the Milky Way and measure the height at which the reversal occurs. Provided this measurement is successful, it would give a measurement of the scale height of the spiral potential. In terms of vertical velocities, within corotation, the mean vertical motion is directed away from the plane at the outer edge of the arms and towards the plane at the inner edge of the arms. The patterns of $\langle v_z \rangle$ above and below the plane are thus mirror-images (see e.g. Carlin et al. 2013). The direction of the mean vertical motion changes roughly in the middle of the interam region. This produces diagonal features in terms of isocontours of a given $\langle v_z \rangle$, as observed by Williams et al. (2013). The picture that emerges from our simulation is one of ``source'' points of the velocity flow in the meridional plane, preferentially on the outer edge of the arms (inside corotation, whilst on the inner edge outside corotation), and of ``sink'' points, preferentially on the inner edge of the arms (inside corotation), towards which the mean velocity flows. We have then shown that this qualitative structure of the mean velocity field is also the behaviour of the analytic solution to linearized Euler equations for a toy model of a cold fluid in response to a spiral perturbation. In a more realistic analytic model, this fluid velocity would in fact be damped by a reduction factor depending on both radial and vertical velocity dispersions when treating the full linearized Boltzmann equation. In a next step, the features found in the present test-particle simulations will also be checked for in fully self-consistent simulations with transient spiral arms, to check whether non-zero mean vertical motions as found here are indeed generic. The response of the gravitational potential itself to these non-zero motions should also have an influence on the long-term evolution of the velocity patterns found here, in the form of e.g. bending and corrugation waves. The effects of multiple spiral patterns (e.g., Quillen et al. 2011) and of the bar (e.g., Monari et al. 2013, 2014) should also have an influence on the global velocity fiel and on its amplitude. Once all these different dynamical effects and their combination will be fully understood, a full quantitative comparison with present and future datasets in 3D will be the next step. The present work on the orbital response of the thin disc to a small spiral perturbation by no means implies that no external perturbation of the Milky Way disc happened in the recent past, by e.g. the Sagittarius dwarf (e.g., Gomez et al. 2013). Such a perturbation could of course be responsible for parts of the velocity structures observed in various recent large spectrosocpic surveys. For instance, concerning the important north-south asymmetry spotted in stellar densities at relatively large heights above the disc, spiral arms are less likely to play an important role. Nevertheless, any external perturbation will also excite a spiral wave, so that understanding the dynamics of spirals is also fundamental to understanding the effects of an external perturber. The qualitative similarity between our simulation (e.g., Fig.~\ref{f:cartevzRZ}), as well as our analytical estimates for the fluid approximation (Fig.~\ref{f:euler}), and the velocity pattern observed by Williams et al. (2013, their Fig.~13) indicates that spiral arms are likely to play a non-negligible role in the observed velocity pattern of our ``wobbly Galaxy''. | 14 | 3 | 1403.0587 |
1403 | 1403.2075_arXiv.txt | {For the first time we explore the circumstellar effects on the Rb (and Zr) abundance determination in O-rich asymptotic giant branch (AGB) stars by considering the presence of a gaseous circumstellar envelope with a radial wind. A modified version of the spectral synthesis code Turbospectrum was used to deal with extended atmosphere models and velocity fields. The Rb and Zr abundances were determined from the resonant 7800 \AA\ Rb I line and the 6474 \AA\ ZrO bandhead, respectively, in five representative O-rich AGB stars with different expansion velocities and metallicities. By using our new dynamical models, the Rb I line profile (photospheric and circumstellar components) is very well reproduced. Interestingly, the derived Rb abundances are much lower (by 1-2 dex) in those O-rich AGB stars showing the higher circumstellar expansion velocities. The Zr abundances, however, remain close to the solar values. The Rb abundances and [Rb/Zr] ratios derived here significantly resolve the problem of the present mismatch between the observations of intermediate-mass (4--8 M$_{\sun}$) Rb-rich AGB stars and the AGB nucleosynthesis theoretical predictions.} | The asymptotic giant branch (AGB) is the last nuclear-burning phase of low- and intermediate-mass stars (0.8 $\leq$ M $\leq$ 8 M$_{\sun}$). The AGB stars are among the main contributors to the chemical enrichment of the interstellar medium since they suffer strong mass loss together with nucleo\-synthesis processes \citep[see e.g.][]{busso1999}. Low-mass AGB stars (M $<$ 4 M$_{\sun}$) can turn C-rich (i.e. C/O $>$ 1) because of the dredge-up of carbon from the bottom of the convective envelope to the ste\-llar surface. The $s$-process a\-llows the production of elements heavier than iron by $slow$ neutron-captures. In these stars, the $^{13}$C($\alpha$, n)$^{16}$O reaction is assumed to ope\-rate as the main neutron source \citep[e.g.][]{abia2001}. On the other hand, intermediate-mass AGB stars ( 4 $\leq$ M $\leq$ 8 M$_{\sun}$) are O-rich stars (C/O $<$ 1) because of the operation of hot bo\-ttom burning (HBB), which burns carbon at the base of the convective envelope, thus preventing the formation of a carbon star \citep{sackmann1992}. In these carbon stars, the $s$-process elements are expected to form mainly by the neutrons released by the $^{22}$Ne($\alpha$, n)$^{25}$Mg reaction, in a higher neutron density environment than in lower mass AGB stars \citep{garcia2006}. The relative abundance of s-elements such as Rb to other neighboring ones such as Sr, Y, and Zr is an indicator of the neutron density, namely a discrimi\-nant of the stellar mass and the main neutron source at the $s$-process site \citep{lambert1995,abia2001,garcia2006}. Observationally, a low [Rb/Zr] ratio ($<$ 0) is found in low-mass AGB stars \citep{plez1993,lambert1995,abia2001} while higher mass AGB stars display [Rb/Zr] $>$ 0 \citep{garcia2006,garcia2007,garcia2009}. \citet[][hereafter Paper I and Paper II, respectively]{garcia2006,garcia2009} derived the Rb and Zr abundances in several Galactic and Magellanic Cloud intermediate-mass AGB stars among a sample of OH/IR stars. The Rb abundances and [Rb/Zr] ratios found in these objects re\-present a challenge for theoretical AGB nucleosynthesis mo\-dels, which do not predict the extreme Rb overabundances ([Rb/Fe] $\gtrsim$ 2 dex) and extraordinarily high [Rb/Zr] ratios observed \citep{vanraai2012,karakas2012}. However, the Rb abundance was derived from the resonant Rb I absorption line at 7800 \AA, using hydrostatic model atmospheres. The Rb I line is probably affec\-ted by contamination from one or more circumstellar (CS) components, as has already been suggested by the detection of blue-shifted CS Rb I absorption lines in several of these extreme O-rich AGB stars (see Paper I, II). In the present {\it Letter}, we explore for the first time the CS effects on the Rb and Zr abundances derived in extreme O-rich AGB stars. To this end, we use more realistic model atmospheres that include a gaseous CS envelope. The much lower Rb abundances (and [Rb/Zr] ratios) derived here significantly alleviate the actual mismatch between the AGB nucleosynthesis predictions and the optical observations of intermediate-mass (4--8 M$_{\sun}$) Rb-rich AGB stars. | The parameters of the dynamical atmosphere models providing the best fit to the observations and the derived Rb and Zr abundances ([Rb/M]$_{dyn}$ and [Zr/M]$_{dyn}$) are shown in Table \ref{table:2}. The Rb abundances ([Rb/M]$^{ref}_{static}$) as determined in Papers I and II from hydrostatic models are also shown for comparison. We note, however, that Papers I and II used solar abundances from \cite{grevesse1998} while we assume here the most recent solar composition by \cite{grevesse2007}. For this reason, Table \ref{table:2} also lists the re-derived Rb abundances ([Rb/M]$_{static}$) using the hydrostatic models with our adopted solar abundances. Our static Rb abundances agree well, within the errors, with those previously derived in Papers I and II. \begin{table*} \tiny \begin{center} \caption{Atmosphere parameters and abundances derived using dynamical models vs. hydrostatic models.} \label{table:2} \centering \begin{tabular}{c c c c c c| c c c c} \noalign{\smallskip} \hline\hline \noalign{\smallskip} IRAS name&$T\rm{_{eff}}$ (K)&$\rm \log$ g& $\beta$& \textit{\.M} (M$_{\sun}$ yr$^{-1}$)&$v$ (km s$^{-1}$)&[Rb/M]$^{ref^a}_{static}$&[Rb/M]$_{static}$$^b$&[Rb/M]$_{dyn}$$^b$&[Zr/M]$_{dyn}$$^b$\\ \hline \noalign{\smallskip} \multicolumn{10}{c}{Galactic stars}\\ \hline \noalign{\smallskip} 05098$-$6422 &3000 & $-$0.5 & 1.0& 1.0 $\times$ 10$^{-8}$ &6 & 0.1 &0.0 $\pm$ 0.4&0.0 $\pm$ 0.4 &$\leq$ 0.3 $\pm$ 0.3 \\ 06300$+$6058 &3000 & $-$0.5 & 0.2& 1.0 $\times$ 10$^{-7}$ &12& 1.6 &1.9 $\pm$ 0.4&0.5 $\pm$ 0.7 &$\leq$ 0.1 $\pm$ 0.3 \\ 18429$-$1721 &3000 & $-$0.5 & 1.0& 1.0 $\times$ 10$^{-8}$ &7& 1.2 &1.2 $\pm$ 0.4&1.0 $\pm$ 0.4 &$\leq$ 0.3 $\pm$ 0.3 \\ 19059$-$2219 &3000 & $-$0.5 & 0.4& 1.0 $\times$ 10$^{-7}$ &13& 2.3/2.6&2.4 $\pm$ 0.4&0.8 $\pm$ 0.7 &$\leq$ 0.3 $\pm$ 0.3 \\ \noalign{\smallskip} \hline \multicolumn{10}{c}{LMC star}\\ \hline\noalign{\smallskip} 04498$-$6842 &3400 & 0.0 & 1.0& 1.0 $\times$ 10$^{-7}$ &13& 3.9$^{c}$ &3.3 $\pm$ 0.4&1.5 $\pm$ 0.7 &$\leq$ 0.3 $\pm$ 0.3 \\ \hline \end{tabular} \tablefoot{ \\ \tablefoottext{a}{See Paper I and II.} \tablefoottext{b}{The uncertainties represent the formal errors due to the sensitivity of the derived abundances to slight changes in the model atmosphere parameters ($\Delta${\it T}$_{eff}$=$\pm$100 K, $\Delta$[M/H]=$\pm$0.3, $\Delta$$\xi$=$\pm$1 kms$^{-1}$, $\Delta$log {\it g}=$+$0.5, $\Delta$FWHM=50 m\AA, $\Delta$$\beta$=0.2, $\Delta$log(\.M/M$_{\odot}$yr$^{-1}$)=1.0}) for each star. \tablefoottext{c}{We scale the Rb overabundance derived by Paper II, [Rb/M] = +5.0, to the adopted LMC metallicity [M/H] = $-$0.3.}} \end{center} \end{table*} \begin{figure} \includegraphics[width=7cm,angle=-90]{fig/RbIpaper.eps} \caption{Rb abundances derived in the sample stars using dynamical models. The location of the Rb I stellar line is indicated by a dashed line. Dynamical models providing the best fits to the observations (black dots) are indicated by a red line. Hydrostatic models are also shown for comparison (blue lines). The expansion velocity and the mass-loss rate adopted in the models also are indicated for each star.} \label{Rb} \end{figure} Interestingly, the new Rb abundances (in the range [Rb/M]$\sim$0.0$-$1.5 dex; Table \ref{table:2}) derived from our dynamical models display a dramatic decrease of 1.4 dex to 1.8 dex with respect to the static case (Table \ref{table:2}) in the more extreme stars, namely those stars showing the highest expansion velocities and mass-loss rates such as IRAS 06300$+$6058, IRAS 19059$-$2219 and IRAS 04498$-$6842. However, for the less extreme stars with a lower expansion velocity and mass-loss rate (IRAS 05098$-$6422 and IRAS 18429$-$1721), the Rb abundances obtained from dynamical models remain close, within 0.2 dex, to those from hydrostatic models. In Fig. \ref{Rb}, we display the observed Rb I line profiles in our O-rich AGB sample (black dots) together with the best synthetic spectra as obtained from the new dynamical models (red lines) versus the static ones (blue lines). Our dynamical atmosphere models reproduce the observed Rb I line profiles (photospheric and circumstellar components) very well, much better than the classical hydrostatic models. The 6474 \AA\ ZrO bandhead, because it is formed deeper in the atmosphere, is less affected than the Rb I line, and the Zr abundances (nearly-solar; see Fig. \ref{Zr}) derived from dynamical models are similar to those obtained with the hydrostatic models (see Papers I and II). Remarkably, we find that relatively low mass-loss rates ($\sim$10$^{-7}$$-$10$^{-8}$ M$_{\odot}$ yr$^{-1}$) give superior fits to the observed Rb I line profiles. Higher mass-loss rates of $\geq$10$^{-6}$ M$_{\odot}$ yr$^{-1}$ give Rb I absorption lines that are too strong even for solar Rb abundances (see Fig. \ref{RbK}). According to \cite{justtanont2013}, the envelope size of $\sim$1x10$^{16}$ cm and the mass-loss rate of $\sim$10$^{-4}$ M$_{\odot}$ yr$^{-1}$ estimated in optically obscured OH/IR AGB stars can be taken as upper limits for massive AGB stars with optical counterparts; i.e. like our sample stars with useful spectra around the Rb I line because the heavily obscured OH/IR massive AGB stars studied by \cite{justtanont2013} have already entered the superwind phase. Our best fits for the Rb I 7800 \AA\ line profiles in optically bright massive AGB stars are those with mass-loss rates lower than $\sim$10$^{-6}$ M$_{\odot}$ yr$^{-1}$. Thus, we conclude that the Rb I line profiles observed in massive AGB stars are consistent with these stars still experiencing relatively low mass-loss rates just before the superwind phase when the Rb density and the envelope size are not very high and this is supported by the MCS presented above (Sect. 3.3). Standard nucleosynthesis models for intermediate-mass AGB stars show that the predicted Rb abundances range from [Rb/M]$\sim$0.0 up to 1.44 dex, depending on the progenitor mass and metallicity \citep[see][]{vanraai2012}; the predicted Rb production increases with increasing stellar mass and decreasing metallicity. Maximum [Rb/M] overabundances of 1.04 and 1.44 are found for a solar metallicity 6.5 M$_{\sun}$ star and for a LMC metallicity 6 M$_{\sun}$ star, respectively \citep{vanraai2012}. More recently, \cite{karakas2012} have delayed the beginning of the superwind phase in solar metallicity nucleosynthesis models of massive AGB stars. These models produce more Rb than in the standard van Raai et al. models because the star experiences more thermal pulses before the superwind phase at the very end of the AGB. The maximum Rb production ([Rb/M]=1.34 dex) is predicted to occur for the 6 M$_{\sun}$ case. By considering the error bars in the spectroscopic analysis (see Table \ref{table:2}) and the theoretical uncertainties \citep[see e.g.][]{vanraai2012,karakas2012}, the Rb abundances are now in fair agreement with the massive AGB nucleosynthesis models, both standard and with delayed superwinds. The nearly-solar derived Zr abundances in IRAS 05098$-$6422, IRAS 06300$+$6058, and IRAS 19059$-$2219 (Table \ref{table:2}) translate into [Rb/Zr] ratios of $-$0.3, 0.4, and 0.5, respectively, which agree quite well with the theoretical predictions (-0.2 $\leq$ [Rb/Zr] $<$0.6). However, the [Rb/Zr] ratios in IRAS 18429-1721 (0.7) and IRAS 04498-6842 (1.2) are still higher than predicted. As already pointed out in the literature \citep[see e.g.][]{vanraai2012}, a possible solution to this observational problem is that gaseous Zr, with a condensation temperature (1741 K) higher than that for Rb (800 K), condensates into dust grains, producing the apparent Zr underabundance that we measure from the ZrO bandheads. In summary, the Rb abundances and [Rb/Zr] ratios derived here significantly resolve the problem of the present mismatch between the observations of massive (4--8 M$_{\sun}$) Rb-rich AGB stars and the theoretical predictions. In the near future, we plan to carry out a chemical analysis based on these new dynamical models for all the Rb-rich AGB stars already studied in Papers I and II. This undoubtedly will help us to constrain the actual nucleosynthesis models for the more massive AGB stars. | 14 | 3 | 1403.2075 |
1403 | 1403.7699_arXiv.txt | We report on {\it Spitzer Space Telescope} IRAC observations of near-Earth object (NEO) 2009~BD that were carried out in support of the NASA Asteroid Robotic Retrieval Mission (ARRM) concept. We did not detect 2009~BD in 25~hrs of integration at 4.5~$\mu$m. Based on an upper-limit flux density determination from our data, we present a probabilistic derivation of the physical properties of this object. The analysis is based on the combination of a thermophysical model with an orbital model accounting for the non-gravitational forces acting upon the body. We find two physically possible solutions. The first solution shows 2009~BD as a $2.9\pm0.3$~m diameter rocky body ($\rho=2.9\pm0.5$~g~cm$^{-3}$) with an extremely high albedo of $0.85_{-0.10}^{+0.20}$ that is covered with regolith-like material, causing it to exhibit a low thermal inertia ($\Gamma=30_{-10}^{+20}$~SI units). The second solution suggests 2009~BD to be a $4\pm1$~m diameter asteroid with $p_V=0.45_{-0.15}^{+0.35}$ that consists of a collection of individual bare rock slabs ($\Gamma = 2000\pm1000$~SI units, $\rho = 1.7_{-0.4}^{+0.7}$~g~cm$^{-3}$). We are unable to rule out either solution based on physical reasoning. 2009~BD is the smallest asteroid for which physical properties have been constrained, in this case using an indirect method and based on a detection limit, providing unique information on the physical properties of objects in the size range smaller than 10~m. | \label{lbl:introduction} The physical properties of near-Earth objects (NEOs) provide important hints on their origin, as well as their past physical and orbital evolution. The most accessible physical properties are the diameter, $d$, and the geometric albedo, $p_V$, which have been measured for more than 1000 NEOs with diameters down to slightly less than 100~m in two large-scale programs, the Warm {\it Spitzer} NEO survey ``ExploreNEOs'' \citep{Trilling2010}, and the ``NEOWISE'' project \citep{Mainzer2011}, using the {\it Wide-field Infrared Survey Explorer} \citep[{\it WISE},][]{Wright2010}. Recently, \citet{Mainzer2013} measured the sizes and albedos of the smallest optically discovered NEOs ($d > 10$~m) from NEOWISE data. Little is known about the physical properties of even smaller NEOs, which constitute the bulk of the NEO population. Knowledge of the physical properties of such small NEOs, some of which pose an impact threat to the Earth, is of importance for understanding their evolution and estimating the potential of destruction in case of an impact, as well as for designing the most promising mitigation mission. Further information on asteroid physical properties are available only for select objects with relatively large sizes, which make up only a fraction of the whole asteroid population. Such properties include, but are not limited to, the bulk density, $\rho$, thermal inertia, $\Gamma$, and the obliquity, $\gamma$, all of which affect non-gravitational forces that act upon the body and alter its orbit compared to a Keplerian one. Two important effects are the Yarkovsky effect \citep[see, e.g.,][]{Bottke2006} and the solar radiation pressure \citep{Vokrouhlicky2000}. The bulk density, $\rho$, provides the simplest way of gaining insight into asteroid interiors. Solid rock, or monolithic, bodies have high bulk densities ($\rho\sim3$~g~cm$^{-3}$), whereas those of rubble-pile bodies, aggregates of smaller particles that are consolidated by their self-gravity or other adhesive forces \citep{Chapman1978}, can be significantly lower as a result of ``macroporosity.'' Macroporosity refers to cavities and void spaces that occur between the irregularly shaped individual constituents \citep[see][for a discussion]{Richardson2002, Britt2002}. \citet{Britt2002} found that most asteroids show a significant degree of macroporosity, in support of the hypothesis that most asteroids must have been disrupted in the course of high-velocity impacts over the age of the Solar System \citep{Chapman1978}. Small asteroids are generally thought of as being individual pieces of compact debris that were generated in disruptive collisions \citep{Pravec2002}; hence, their macroporosity is expected to be low and their bulk density high compared to that of rubble pile asteroids. Thermal inertia, $\Gamma$, describes the ability of the surface material to store thermal energy: high-thermal-inertia material heats up slowly and re-emits the thermal energy only gradually, whereas low-thermal-inertia material can be approximated as being in instantaneous thermal equilibrium with the incoming insolation \citep[see, e.g.,][]{Spencer1989}. Examples for materials of low and high thermal inertia are regolith \citep[30--50~SI units,][1~SI unit equals 1~J~m$^{-2}$~s$^{-0.5}$~K$^{-1}$]{Spencer1989,Putzig2005} and bare rock \citep[${>}2500$~SI units,][]{Jakosky1986}, respectively. Measurements of the thermal inertia of medium-to-large sized NEOs ($d > 100$~m) revealed values of 100--1000~SI units \citep{Delbo2007}. Both the thermal inertia and the bulk density of asteroids can be derived by modeling the effect of non-gravitational perturbations on the object's orbit \citep[see, e.g.,][]{Chesley2014}. Assuming a homogeneous bulk density of the constituent particles, usually derived from laboratory measurements of meteorite equivalent material, allows for constraining the degree of macroporosity of the asteroid. NEO 2009~BD was discovered on January 16, 2009, at a distance from the Earth of only 0.008~au \citep{Buzzi2009}. Its orbit is very Earth-like with a period of 400~days (JPL Solution 41). The escape velocity of 2009~BD with respect to the Earth is among the lowest for known objects ($v_{\infty} \sim 1$~km~s$^{-1}$), making it a worthwhile candidate mission target. 2009~BD is considered the primary candidate mission target for NASA's Asteroid Robotic Retrieval Mission \citep[ARRM,][]{NASAAsteroidInitiativeWebsite}. The mission concept involves capturing an asteroid and dragging it onto a new trajectory that traps it in the Earth-Moon system, where it will be further investigated by astronauts. As a result of 2009~BD's Earth-like orbit, its next encounter with the Earth-Moon system will be in late 2022, when the proposed capture through ARRM would take place. The current mission design requires the target asteroid to have a diameter of 7--10~m and a total mass of ${\sim}500$~metric tons \citep{NASAmissionwebsite}. The orbital parameters and absolute magnitude, $H$, which is the apparent magnitude of an object at a distance of 1~au to the Sun and the observer, of 2009~BD are well-known, providing accurate orbital predictions \citep{Micheli2012}. However, there is no albedo-independent determination of its diameter, which is a crucial variable in the ARRM mission planning. We report here in observations of 2009~BD using the IRAC camera on the {\it Spitzer Space Telescope}, which provides the only practical means to constrain the physical properties of 2009~BD in the next decade. The main goals of our observations were two-fold: measure the size and therefore determine the suitability of 2009~BD as an ARRM mission target, and constrain other physical properties like bulk density and thermal inertia of an asteroid at a size range that is so far unprecedented. | \label{lbl:discussion} \subsection{Observations and {\it Spitzer} Pointing} \label{lbl:discussion_observations} 2009~BD was not detected in our observations, and was fainter than expected (see Figure \ref{fig:sensitivity_finalimage}). Furthermore, the final sensitivity was lower than predicted: a 5$\sigma$ detection would have been anticipated at $0.78$~$\mu$Jy in channel 2, whereas we are only able to derive a $3\sigma$ upper limit at this flux density. The lower sensitivity is possibly a result of the fact that 2009~BD moved only a short distance over the integration time, compromising the background removal and increasing the background noise of the final mosaic. To verify that our non-detection is meaningful we have carefully checked to ensure that the {\it Spitzer} pointing was correct and the 2009~BD ephemeris predictions were accurate. To check the {\it Spitzer} pointing we have compared the most current predicted position with the position assumed by {\it Spitzer} for each frame. We find a mean discrepancy between the {\it Spitzer} pointing and the orbital predictions of $4.7\pm0.3$\arcsec\ in right ascension and $1.6\pm0.2$\arcsec\ in declination. This small discrepancy is readily explained by the fact that the predictions are based on a more recent {\it Spitzer} ephemeris than was available during the {\it Spitzer} observations; {\it Spitzer} position errors can be a few hundred km, leading to 3-6\arcsec\ pointing errors for 2009~BD. Therefore, the new position prediction is superior to the one assumed at the time of the observations. The ellipse in Figure \ref{fig:sensitivity_finalimage} marks this updated position, which does not show a detection of 2009~BD. So far we have shown that the searched position matches the predicted position, but to have confidence in the non-detection analysis we must also show that 2009~BD was in fact close to the predicted position. To this end we compared our nominal plane-of-sky predictions from JPL Solution 41 to the results of various alternate orbital solutions, e.g., taking into account $A1$ only, both $A1$ and $A2$, gravitational effects only, and different outlier rejection schemes \citep{Carpino2003}, with the result that 2009~BD was always near the nominal prediction. We also tried to fit synthetic observations that were ${\sim}5\arcsec$ from the prediction, but this consistently led to unrealistic residuals for the other observations. From this we conclude that 2009~BD could not be even ${\sim}5\arcsec$ from its predicted position. \subsection{Physical Properties} \label{lbl:discussion_properties} By combining the orbital and the thermophysical model we are able to constrain a variety of physical properties of 2009~BD by taking advantage of the mutual dependencies of the individual properties. The range of each physical property is confined based on purely physical considerations. The only properties that are not derived as part of the modeling process are $H$, $G$, and $P$ (see Section \ref{lbl:modeling}), which are based on observations of 2009~BD or general properties of the asteroid population. We discuss those physical properties that are constrained in the modeling process. We derive a volume-equivalent diameter of $2.9\pm0.3$~m and a geometric albedo of $p_V = 0.85_{-0.10}^{+0.20}$ for the low-$\Gamma$ solution and $d=4\pm1$~m and $0.45_{-0.15}^{+0.35}$ for the high-$\Gamma$ solution. Both albedo solutions exhibit relatively high albedos compared to albedo measurements of other NEOs \citep{Trilling2010, Mainzer2011}. The diameter results agree within uncertainties, allowing for a more generalized formulation: 2009~BD's diameter is $2.6 < d < 7$~m at a 3$\sigma$ confidence level. Note that despite the fact that we did not detect 2009~BD, we are able to constrain the diameter in both cases with uncertainties that are rather low compared to the typical 20\% diameter uncertainty derived from ExploreNEOs data \citep{Harris2011}. We are able to confine the diameter to such a narrow range due to constraints given by the physics of the orbital model and the upper-limit flux density derived from our observations. As a matter of fact, a diameter lower than 2.6~m implies a high Bond albedo, which in turn reduces the size of the Yarkovsky effect and prevents our model from matching the magnitude of the observed acceleration. The advantage of this definition of the lower limit is that we do not {\it apriori} rule out the possibility that 2009~BD has an extremely high albedo, which is the case for the low-$\Gamma$ solution. Although relatively rare, similarly high albedos have been found by both the ExploreNEOs \citep{Trilling2010} and the NEOWISE projects \citep{Mainzer2011}. Based on spectral work by \citet{Thomas2011}, both albedo results are compatible with a E/S/V/Q-type taxonomic classification for 2009~BD. We compare the derived bulk density solutions for 2009~BD with measured bulk densities of asteroids with diameters of 10~km or smaller. Density data are taken from the list compiled by \citet{Baer2012} and the literature (see Table \ref{tbl:physical_properties}). All measurements are plotted in Figure \ref{fig:physical_properties}. Macroporosity is derived as unity minus the ratio of the asteroid's bulk density and the bulk density of meteorite equivalent material. Meteorite equivalent bulk densities are taken from \citet{Britt2002}, Table 2. Due to ambiguities in the identification of the asteroidal origin of meteoritic material, we use average bulk densities derived from meteoritic material as proxies for asteroidal bulk densities of individual taxonomic types. Hence, we assume S and Q-type asteroids to have an average bulk density of $3.3\pm0.1$~g~cm$^{-3}$, as derived from H/L/LL ordinary chondrites, and C-type and B-type asteroids to have an average bulk density of $2.6\pm0.5$~g~cm$^{-3}$, as derived from different types of carbonaceous chondrites. For V-type asteroids we assume an average bulk density of $2.9_{-0.4}^{+0.5}$~g~cm$^{-3}$ as derived from howardite-eucrite-diogenite meteorites \citep{Macke2011}. \citet{Consolmagno2008} give a mean bulk density of enstatite chondrites, which likely originate from E-type asteroids, of $3.5\pm0.2$~g~cm$^{-3}$. Figure \ref{fig:physical_properties} (center panel) shows that most objects with diameters 100~m $< d <$ 10~km have macro-porosities higher than 30\%, consistent with a rubble-pile nature \citep{Britt2002}. There is no clear trend in either bulk density or macroporosity with the diameter of the object. Using the derived bulk density of 2009~BD ($2.9_{-0.5}^{+0.5}$~g~cm$^{-3}$ for the low-$\Gamma$ solution, $1.7_{-0.4}^{+0.7}$~g~cm$^{-3}$ for the high-$\Gamma$ solution), and a mean bulk density of $(3.2\pm0.3)$~g~cm$^{-3}$ derived as the average of the S/Q/V/E-type asteroid material bulk densities listed above, we derive a degree of macroporosity of $10_{-10}^{+20}$\% or $45_{-30}^{+15}$\%. Hence, our low-$\Gamma$ solution is most consistent with a monolithic nature of 2009~BD, whereas the high-$\Gamma$ solution suggests a rubble-pile nature \citep{Britt2002}. \begin{deluxetable}{rccccccl} \tabletypesize{\scriptsize} \tablecaption{Physical Properties of Small Asteroids\label{tbl:physical_properties}} \tablewidth{0pt} \tablehead{ \colhead{Object} & \colhead{Diameter} & \colhead{Albedo} & \colhead{Tax.} & \colhead{Bulk Density} & \colhead{MPor.} & \colhead{$\Gamma$} & \colhead{Ref.} \\ & \colhead{(km)} & & Type & \colhead{(g cm$^{-3}$)} & \colhead{(\%)} & \colhead{(SI)} & } \startdata (1580)\ Betulia & $4.57\pm0.46$ & $0.08\pm0.02$ & C & ... & ... & $180\pm50$ & 1\\ (1862)\ Apollo & $1.55\pm0.07$ & $0.20\pm0.02$& Q & $2.85\pm0.68$ & $15_{-15}^{+25}$ & $140_{-100}^{+140}$ & 2, 3\\ (3749)\ Balam & $(7.2\pm0.4)$\tablenotemark{a} & (0.15)\tablenotemark{a} & (S)\tablenotemark{b} & ($2.61\pm0.45$)\tablenotemark{a} & $20_{-15}^{+15}$& ... & 4 \\ (3908)\ Nyx & $1.0\pm0.2$ & $0.15\pm0.08$ & V & $0.9\pm0.2$\tablenotemark{c} & $70_{-15}^{+10}$ & ... & 5, 6, 7 \\ (25143)\ Itokawa & $0.320\pm0.001$ & $0.30\pm0.10$ & S & $1.90\pm0.13$ & $40_{-10}^{+10}$ & $700\pm100$ & 1, 8, 9 \\ (33342)\ 1998~WT24 & $0.35\pm0.04$ & $0.6\pm0.2$ & E & ... & ... & $200\pm100$ & 1, 10\\ (54509)\ YORP & $0.09\pm0.01$ & $0.20\pm0.02$ & S/V & ... & ... & $700\pm500$ & 1 \\ (66391)\ 1999~KW4 & $1.33\pm0.07$ & ($0.25$)\tablenotemark{d}& S & $2.00\pm0.26$ & $40_{-10}^{+10}$ & ... & 5, 11\\ (101955)\ Bennu & $0.50\pm0.02$ & $0.05\pm0.01$ & B & $1.2\pm0.1$ & $55_{-15}^{+10}$ & $650\pm100$ & 12, 13\\ (162173)\ 1999~JU3 & $0.87\pm0.03$ & $0.07\pm0.01$ & C & ... & ... & $400\pm200$ & 3, 14\\ (175706)\ 1996~FG3 & $1.71\pm0.07$ & $0.044\pm0.004$ & C & ... & ... & $120\pm50$ & 15, 16\\ (185851)\ 2000~DP107 & $0.81\pm0.18$ & ($0.14$)\tablenotemark{d} & (S)\tablenotemark{d} & $1.65\pm0.84$ & $50_{-30}^{+25}$ & ... & 17\\ (308635)\ 2005~YU55 & $0.31\pm0.01$ & $0.07\pm0.01$ & C & ... & ... & $580\pm230$ & 18, 19\\ (341843)\ 2008~EV5 & $0.37\pm0.01$ & $0.13\pm0.05$ & C & ... & ... & $450\pm60$ & 19, 20 \\ 2000~UG11 & $0.23\pm0.03$ & (0.23)\tablenotemark{d} & (S)\tablenotemark{d} & $1.47\pm0.7$ & $55_{-25}^{+20}$ & ... & 21\\ 2002~CE26 & $3.50\pm0.40$ & 0.07 & (C)\tablenotemark{e} & $0.9\pm0.5$ & $65_{-30}^{+20}$ & ... & 22\\ 2002~NY40 & $0.28\pm0.03$ & $0.34\pm0.06$ & Q & ... & ... & $550\pm450$ & 23, 24 \\ 2003~YT1 & $1.06\pm0.06$ & (0.52)\tablenotemark{d} & V & $2.01\pm0.70$ & $30_{-30}^{+30}$ & ... & 25, 26\\ \enddata \tablecomments{This table lists measured physical properties of asteroids with diameters of 10~km or smaller. Data in brackets are based on assumptions (see below). In the case of multi-component systems, diameters and bulk densities refer to the average numbers of the combined system (diameters are those of a volume-equivalent sphere). The macroporosity (``MPor.'') is derived through division of the asteroid's bulk density by the bulk density of the respective meteorite equivalent material (see text for details). Comparison data for 2009~BD can be found in Table \ref{tbl:solutions}.} \tablenotetext{a}{based on an assumed albedo} \tablenotetext{b}{based on its Flora family membership \citep{Marchis2008}} \tablenotetext{c}{density estimate assumes a thermal inertia according to \citet{Delbo2007}} \tablenotetext{d}{albedo derived from the equivalent diameter and the $H$ magnitude \citep{SBDB}} \tablenotetext{e}{taxonomic type assigned based on albedo determinations by \citet{Thomas2011}} \tablerefs{ (1) \citet{Mueller2007}; (2) \citet{Rozitis2013}; (3) \citet{Bus2002}; (4) \citet{Marchis2008}; (5) \citet{Binzel2004}; (6) \citet{Benner2002}; (7) \citet{Farnocchia2014b}; (8) \citet{Fujiwara2006}; (9) \citet{Abe2006}; (10) \citet{Kiselev2002}; (11) \citet{Ostro2006}; (12) \citet{Mueller2012}; (13) \citet{Chesley2014}; (14) \citet{Mueller2011b}; (15) \citet{Wolters2011}; (16) \citet{Thomas2011}; (17) \citet{Margot2002}; (18) \citet{Mueller2013}; (19) \citet{Somers2010}; (20) \citet{Ali-Lagoa2013}; (21) \citet{Margot2002b}; (22) \citet{Shepard2006}; (23) \citet{Roberts2007}; (24) \citet{Mueller2004}; (25) \citet{Brooks2006}; (26) \citet{Sanchez2013}.} \end{deluxetable} \begin{figure} \epsscale{.80} \plotone{fig7.eps} \caption{Physical properties of known asteroids with diameters of 10~km or less, as a function of their diameter: bulk density (left), macroporosity (center), and thermal inertia (right). Grey circles depict stony asteroid types (S/Q/V/E), black circles carbonaceous types (B/C) (according to Table \ref{tbl:physical_properties}). The low-$\Gamma$ solution of 2009~BD is indicated as a red diamond, the high-$\Gamma$ solution as a blue diamond in each plot. A large degree of variety in bulk densities and macroporosity is obvious, and there is no clear trend between either of them and diameter. Thermal inertia exhibits a slight trend of smaller objects having higher thermal inertia \citep[see][]{Delbo2007}. Data plotted here are tabulated in Tables \ref{tbl:solutions} and \ref{tbl:physical_properties}.\label{fig:physical_properties}} \end{figure} Previous measurements of the thermal inertia of asteroids revealed values in the range ${\sim}$10--100~SI units for large main belt asteroids and higher values up to 1000~SI units for NEOs \citep[see, e.g.,][and references therein]{Delbo2007}. The lower thermal inertia of large bodies is generally ascribed to the presence of a thick layer of regolith, which has a low thermal inertia \citep[see Section \ref{lbl:introduction}, as well as][]{Spencer1989,Putzig2005}. The right-hand panel of Figure \ref{fig:physical_properties} plots thermal inertia measurements of asteroids with ${\sim}0.1 < d < 10$~km. The plot reveals that 2009~BD's thermal inertia is extreme in this range, irrespective of which of our solutions better describes reality. The low-$\Gamma$ solution is consistent with the presence of regolith, whereas the high-$\Gamma$ solution is consistent with a bare-rock nature of 2009~BD. A slight trend of increasing thermal inertia with decreasing diameter is visible in Figure \ref{fig:physical_properties}, which was already discussed by \citet{Delbo2007}. Our high-$\Gamma$ solution seems to be more consistent with this trend, presuming that this trend is valid for asteroids in the size regime of 2009~BD. We summarize the properties of 2009~BD created by our two solutions in Figure \ref{fig:schematic}. The low-$\Gamma$ solution suggests that 2009~BD is a ${\sim}3$~m sized rocky body ($\rho=2.9$~g~cm$^{-3}$) that is covered with a physically thin but optically thick layer of regolith-like material, causing it to exhibit a low thermal inertia ($\Gamma=30$~SI units) and a high bulk density. This picture seems realistic in the sense that models show that even small asteroids can retain a layer of fine-grained dust on their surfaces \citep{Scheeres2010,Sanchez2013b}. The picture created by the high-$\Gamma$ solutions shows 2009~BD as a 4~m-sized rubble-pile asteroid ($\rho = 1.7$~g~cm$^{-3}$) that consists of individual bare rock slabs ($\Gamma = 2000$~SI units) and exhibits a macroporosity of 45\%. This scenario seems to be more realistic due to the lower albedo that is required for this configuration. Despite the fact that some properties seem to favor one of the two results at a time, we are unable to rule out either of the configurations based on physical reasoning. In order to be able to commit to either solution, additional observations are necessary that are able to pinpoint one decisive physical property of 2009~BD, e.g., its thermal inertia or its bulk density. The thermal inertia can be further constrained using additional infrared observations or using in-situ measurements. The bulk density can be independently derived by measuring the gravitational attraction of the object on a nearby body or a rendezvous spacecraft. \subsection{Discussion of the Modeling Technique} \label{lbl:discussion_modeling} The probabilistic approach taken in this work to constrain the physical properties of 2009~BD is unique. We have to acknowledge that a full validation of the methods presented in this paper is not yet possible. 2009~BD is currently the only asteroid for which both the Yarkovsky and solar radiation pressure forces can be measured from astrometric observations, which are used here to constrain the object's bulk density and thermal inertia. An independent validation of our modeling approach would require an asteroid with non-gravitational perturbations and physical model independently characterized. The first object for which such a wealth of data is anticipated will be NEO (101955) Bennu, the OSIRIS-REx mission target. Instead, we note that both the thermophysical and the dynamical models are individually well-tested. The thermophysical model used in this work is based on and has been extensively tested against the model discussed by \citet{Mueller2007}, which was applied in a number of publications \citep[e.g.,][]{Harris2007, Mueller2010}. The model of the Yarkovsky forces is based on work done by \citet{Vokrouhlicky2000b}, which is used to describe the Yarkovsky effect observed in a number of objects \citep[e.g.,][]{Chesley2003, Vokrouhlicky2008, Farnocchia2013, Chesley2014, Farnocchia2014, Farnocchia2014b}. Also, the solar radiation pressure model \citep{Vokrouhlicky2000} was used to refine orbits of small asteroids \citep[e.g.,][]{Micheli2012, Micheli2012b, Micheli2013}. Note that in all previous works in which either model has been used, the resulting physical properties are within reasonable ranges. Both the orbital and the thermophysical model assume a spherical shape of 2009~BD. \citet{Emery2013} have shown that using the real shape of (101955) Bennu, instead of assuming a spherical shape, lowers the thermal inertia of that object by a factor of 2. The case of Bennu shows that shape information can impact the physical parameter results. We investigate the possible impact of an irregular shape on our results. 2009~BD has a rotation period $P \geq 3$~hrs \citep{Tholen2013} with a lightcurve magnitude of ${\geq}0.25$~mag (B. Ryan, private communication 2014), which suggests an elongation $b/a \leq 0.8$ for a triaxial ellipsoid with relative dimensions ($a$, $b$, $c$). Assuming a rotation period $3 < P \ll 25$~hrs, any elongation effects are averaged out during our 25~hrs integration, leading to physical properties of a volume-equivalent sphere. We further investigate the effect of a possible flattening of 2009~BD, by assuming the shape of a triaxial ellipsoid with axes $b/a = 0.7$ and $c/b = 0.7$, which is quite typical among larger asteroids. We find that the smaller cross-section of the triaxial shape compared to that of a spherical shape requires a reduction of the bulk density of 15--20\% to provide the observed magnitude of the solar radiation pressure force. A numerical simulation of the Yarkovsky forces \citep{Vokrouhlicky2000b} suggests that the thermal inertia estimates might be lower by as much as a factor of 2, which has also been found for Bennu by \citet{Emery2013}. The smaller thermal inertia in turn reduces the diameters found for both solutions (see Figures \ref{fig:d_TI_map} and \ref{fig:diameter_distributions}). Note that the changes to the individual physical properties found as part of this simulation are mostly within the uncertainties derived assuming a spherical shape. Since there is no certainty on the shape of 2009~BD we stick to the results based on a simple spherical shape, which are still valid within the uncertainties, assuming a triaxial shape. Additional information on the shape of 2009~BD might require a re-assessment of our results in the future. \begin{figure} \epsscale{.80} \plotone{fig8.eps} \caption{Schematic diagram of the nature of 2009~BD. Grey lines indicate the different possible configurations as a function of thermal inertia and macroporosity. The red and the blue cloud symbolize the 1$\sigma$ confidence intervals of the low and high-$\Gamma$ solutions, respectively. Note that the border lines of the different configurations are not as strict as shown here. \label{fig:schematic}} \end{figure} \subsection{Implications} Our results show that the volume-equivalent diameter of 2009~BD, $2.6 < d < 7.0$~m (3$\sigma$), is most likely smaller than the size range aimed for in the ARRM mission design (7--10~m, see above). However, its total mass is roughly 1/10 of the mass aimed for in the current design, reducing efforts necessary to alter the orbit of the asteroid. The final decision on 2009~BD's suitability as a mission target is beyond the scope of this work, but a potential mission to 2009~BD will be able to resolve the solution degeneracy we found for the physical properties of this object. The two scenarios based on our data and presented in Section \ref{lbl:discussion_properties} show 2009~BD either as a rocky object covered with regolith-like material or a loose conglomerate of bare rocks. Either scenario reveals this object as rather exotic compared to other known asteroids. Hence, 2009~BD may not belong to the normal population of NEOs that have their origins in the main belt, accounting for its very Earth-like orbit. It has been suggested that ejecta from impacts on the Moon could end up in Earth-like orbits; another possibility is that 2009~BD is a man-made object \citep{Micheli2012}. A value for $p_V$ of around 0.45 or 0.85 is much higher than the Moon's albedo albedo of 0.11 \citep{dePater2001}, which would appear to reduce the likelihood that 2009~BD has a lunar origin. A section of a spent rocket booster would have a high albedo, but the densities derived here (for both the low and the high-$\Gamma$ solutions) are far higher than values associated with, for example, hollow rocket fuel tanks. Specifically, the values of the area-to-mass ratio ($\Psi$) listed in Table \ref{tbl:solutions}, which are substantially independent of the diameter estimate, are 1--2 orders of magnitude less than that of artificial objects. The available astrometry contradicts area-to-mass ratios compatible with a spent booster. While our results do not appear to favor any of the more exotic origins for 2009~BD suggested by its very Earth-like orbit, they emphasize the puzzling nature of this object and the need for further observations of this and similar objects. | 14 | 3 | 1403.7699 |
1403 | 1403.7520_arXiv.txt | The Gemini Planet Imager (GPI) is a dedicated facility for directly imaging and spectroscopically characterizing extrasolar planets. It combines a very high-order adaptive optics system, a diffraction-suppressing coronagraph, and an integral field spectrograph with low spectral resolution but high spatial resolution. Every aspect of GPI has been tuned for maximum sensitivity to faint planets near bright stars. During first light observations, we achieved an estimated $H$ band Strehl ratio of 0.89 and a 5-$\sigma$ contrast of $10^6$ at 0.75 arcseconds and $10^5$ at 0.35 arcseconds. Observations of Beta Pictoris clearly detect the planet, Beta Pictoris b, in a single 60-second exposure with minimal post-processing. Beta Pictoris b is observed at a separation of $434 \pm 6$~ milli-arcseconds (mas) and position angle $211.8 \pm 0.5^\circ$. Fitting the Keplerian orbit of Beta Pic b using the new position together with previous astrometry gives a factor of three improvement in most parameters over previous solutions. The planet orbits at a semi-major axis of $9.0^{+0.8}_{-0.4}$~AU near the 3:2 resonance with the previously-known 6 AU asteroidal belt and is aligned with the inner warped disk. The observations give a 4\% posterior probability of a transit of the planet in late 2017. | Direct detection---spatially resolving the light of a planet from the light of its parent star---is an important technique for characterizing exoplanets. It allows observations of giant exoplanets planets in locations like those in our solar system, which are inaccessible to other methods. The Gemini Planet Imager (GPI) is a new instrument for the Gemini South telescope. Designed and optimized only for high-contrast imaging, it incorporates advanced adaptive optics, diffraction control, and a near-infrared spectrograph and an imaging polarimeter. During first light scientific observations in November 2013, GPI achieved contrast performance that is an order of magnitude better than conventional adaptive optics imagers. \keywords{Adaptive Optics | extrasolar planets | astronomical instrumentation} \dropcap{D}irect imaging is a powerful complement to indirect exoplanet detection techniques. In direct imaging, the planet is spatially resolved from its star, allowing it to be independently studied. This capability opens up new regions of parameter space, including sensitivity to planets at $>$ 5 AU. It also allows spectroscopic analysis of the light emitted or reflected by the planet to determine its composition\cite{Konopacky2013}\cite{Oppenheimer2013} and astrometry to determine the full Keplerian orbital elements\cite{Chauvin2012}\cite{Kalas2013}. Imaging planets is extremely challenging---Jupiter is $10^9$ times fainter than our sun in reflected visible light. Younger extrasolar planets are more favorable targets. During their formation, planets are heated by the release of gravitational potential energy. Depending on the exact formation process and initial conditions, a 4 Jupiter-mass ($M_J$) planet at an age of 10 million years could have a luminosity between $10^{-6}$ and $2 \times 10^{-5} L_\odot$\cite{Marley2007}---but this is still a formidable contrast ratio. To overcome this, astronomers have combined large telescopes (to reduce the impact of diffraction), adaptive optics (to correct for phase errors induced by atmospheric turbulence), and sophisticated image processing \cite{Lafreniere2007}\cite{Soummer2012}. This recipe in various combinations had achieved several notable successes \cite{Kalas2008}\cite{Marois2008}\cite{Lagrange2010}\cite{Rameau2013}\cite{Kuzuhara2013}. However, the rate of these discoveries remains low\cite{Biller2013}\cite{Nielsen2013}\cite{Wahhaj2013}, in part because the number of suitable young stars in the solar neighborhood is low, and, for all but the closest stars, detection is limited to $>20$ AU where planets may be relatively rare. To move beyond this limited sample, dedicated instruments are needed that are designed specifically for high-contrast imaging. One such instrument is the Gemini Planet Imager (GPI). GPI is a fully-optimized high contrast AO facility deployed on an 8-m telescope and is almost an order of magnitude more sensitive than the instruments of the previous generation. With this powerful dedicated facility, large-scale surveys could increase the sample of directly imaged giant planets to 25--50, or more \cite{McBride2011}. We present an overview of the design of GPI, discuss its initial operation and performance, and show first-light science results. | 14 | 3 | 1403.7520 |
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1403 | 1403.2305_arXiv.txt | The absence of thermal instability in the high/soft state of black hole X-ray binaries, in disagreement with the standard thin disk theory, is a long-standing riddle for theoretical astronomers. We have tried to resolve this question by studying the thermal stability of a thin disk with magnetically driven winds in the $\dot{M}- \Sigma$ plane. It is found that disk winds can greatly decrease the disk temperature and thus help the disk become more stable at a given accretion rate. The critical accretion rate $\dot{M}_{\rm crit}$ corresponding to the thermal instability threshold is increased significantly in the presence of disk winds. For $\alpha=0.01$ and $B_{\rm \phi}=10B_{\rm _p}$, the disk is quite stable even for a very weak initial poloidal magnetic field [$\beta_{\rm p,0}\sim 2000, \beta_{\rm p}=(P_{\rm {gas}}+P_{\rm rad})/(B_{\rm p}^2/8\pi)$]. But when $B_{\rm \phi}=B_{\rm _p}$ or $B_{\rm \phi}=0.1B_{\rm _p}$, a somewhat stronger (but still weak) field ($\beta_{\rm p,0}\sim 200$ or $\beta_{\rm p,0}\sim 20$) is required to make the disk stable. Nevertheless, despite the great increase of $\dot{M}_{\rm crit}$, the luminosity threshold corresponding to instability remains almost constant or even decreases slowly with increasing $\dot{M}_{\rm crit}$ due to the decrease of gas temperature. The advection and diffusion timescales of the large-scale magnetic field threading the disk are also investigated in this work. We find that the advection timescale can be smaller than the diffusion timescale in a disk with winds, because the disk winds take away most of the gravitational energy released in the disk, resulting in the decrease of the magnetic diffusivity $\eta$ and the increase of the diffusion timescale. | According to the standard thin disk theory, the radiation pressure dominated inner region of a thin disk is both thermally and viscously unstable when the Eddington-scaled mass accretion rate is larger than a critical value \citep{s1973,l1974,s1976,p1978}, which corresponds to a few percent of the Eddington luminosity. However, the high/soft state of X-ray binaries appears quite stable on observation. \citet{g2004} found that black hole X-ray binaries with luminosities ranging from $0.01$ to $0.5$ $L_{\rm Edd}$ show little variability, which obviously conflicts with the accretion disk theory. Only one superluminous X-ray binary, GRS 1915+105, was found to possess the limit-cycle light curve expected to be produced by thermal-viscous instability over the course of decades \citep{b1997}. The variability of GRS 1915+105 was inferred to be related to its high luminosity \citep{g2004}. But recently, \citet{a2011} reported another source, IGR J17091-3624, that seems to show variability similar to that of GRS 1915+105 at lower luminosity, which suggests that there may be other variables associated with limit-cycle behavior in X-ray binaries. Although the analogous limit-cycle in active galactic nuclei (AGN) is hard to observe directly due to its long timescale, some intermittent activity in young radio galaxies has been ascribed to thermal instability in the disk \citep{c2009,w2009}. There are mainly two processes that can change the theoretical results. Firstly, if the disk viscous stress is proportional to the gas pressure instead of the total pressure, the disk will be stable \citep{s1981,s1984}. But shearing box radiation-MHD simulations by \citet{h2009a} suggested that the stress scales approximately with the total pressure. Simultaneously, the Lightman-Eardley viscous instability was also confirmed. The second method to eliminate the instability is to make the disk cooler, thus increasing the relative importance of gas pressure compared to radiation pressure. \citet{s1994} found that the disk would be stable if most of the gravitational energy released in the disk were transported to the corona. Convective cooling has been suggested as a stabilizing factor \citep{g1995}, although later research showed that it probably has a minor effect on disk stability. Turbulence, instead of convection, has also been suggested to play a key role in increasing the critical accretion rate \citep{z2013}. Another possible mechanism to cool the disk relies on magnetic pressure to provide part of the vertical hydrodynamical support \citep{z2011}. \citet{h2009b} pointed out that the time delay between the turbulent stress and total pressure of the disk can also make the disk stable. But \citet{j2013}, using the same code, found that the disk still runs away once they adopt a large enough horizontal shearing box size. In this work, we investigate the thermal stability of a thin disk with winds. Strong winds driven by a large-scale magnetic field can take away most of the gravitational energy released in the disk, thereby reducing the disk temperature considerably \citep{l2012}. Thus, disk winds can help to cool the disk and make the disk stable. In this work we assume the existence of a large-scale magnetic field threading the disk; how this field is established remains an open question. The formation of large-scale field in a thin disk seems to be difficult due to its fast diffusive speed \citep{v1989,l1994}. \citet{c2013} suggested that the advection timescale can become smaller than the diffusion timescale in the presence of winds, thus the field can be effectively dragged inwards from the outer region even for a thin disk. We consider the advection and diffusion time-scales of the magnetic field based on this work in Section \ref{fields}. | \label{conclusionS} In this work, we investigate the $\dot{M}-\Sigma$ curves of a thin disk with magnetically driven winds. It is found that, because disk winds can greatly decrease the disk temperature, the critical accretion rate $\dot{M}_{\rm crit}$ can be increased significantly and the disk becomes more stable (Figs. \ref{f2}, \ref{f3}). It seems that both the gas and radiation pressure dominated regions possess the same slopes in the $\dot{M}-\Sigma$ curves ($\dot{M}\sim\Sigma^{2\epsilon}$) when magnetic torques drive the inflow. But the real slope in the radiation pressure dominated region is between $-1$ and $2\epsilon$ in numerical calculations because both the viscous torque and magnetic torque are important when the mass accretion rate is close to the Eddington accretion rate. The parameter $\alpha$, the strength and the morphology of the initial magnetic fields all strongly affect the critical accretion rate $\dot{M}_{\rm crit}$. If $\beta_{\rm p,0}$ and $\alpha$ are smaller, the thin disk will be more stable. While the accretion disk with winds becomes stable for a high accretion rate, the luminosity threshold may not increase because of the much lower gas temperature. Indeed, it is found that the disk luminosity corresponding to $\dot{M}_{\rm crit}$ remains almost constant or even decreases slowly with the increase of $\dot{M}_{\rm crit}$ (Fig. \ref{luminosity}). Thus the absence of thermal instability in luminous accretion systems is still a problem even if the disk is stable for a very high accretion rate, unless other components, such as a corona or winds, contribute significantly to the luminosity in relevant spectral bands. Using equation (\ref{betap}) ($\beta_{\rm p}\sim \Sigma^{4/3-2\epsilon}$), it is interesting to note that there is a critical initial field strength $\beta_{\rm p,0,crit}$ corresponding to $\epsilon=2/3$ ($\beta_{\rm p}\sim \Sigma^{0}$). If $\beta_{\rm p,0} < \beta_{\rm p,0,crit}$, $\beta_{\rm p}$ will become smaller and smaller with increasing $\Sigma$ and the disk will tend to be more stable. Otherwise, $\beta_{\rm p}$ will become larger and the disk will be like the standard thin disk with increasing $\Sigma$ if $\beta_{\rm p,0} > \beta_{\rm p,0,crit}$. Thus, the slopes in the $\dot{M}-\Sigma$ curves will tend to be either larger or smaller with increasing $\Sigma$ (see Figs. \ref{f2}, \ref{f3}). We have studied how $\epsilon$ varies with the surface density $\Sigma$ for different initial poloidal fields at $R=2R_{\rm ISCO}$, for example, in Fig. \ref{f10}, where $\alpha=0.01$ and $B_{\phi}=10B_{\rm p}$ are adopted. Only the dash-dotted line ($\beta_{\rm p,0}=2000$) satisfies $\beta_{\rm p,0} < \beta_{\rm p,0,crit}$ for the surface densities considered. Furthermore, both the field strength and $\epsilon$ increase with increasing $\Sigma$. On the contrary, $\beta_{\rm p}$ becomes larger with increasing $\Sigma$ for both the dashed and dotted lines at first, because $\beta_{\rm p,0} > \beta_{\rm p,0,crit}$. As a result, the magnetic field will be unimportant and $\epsilon$ is almost the same as that of a standard disk (the black line). While the disk seems to be quite stable in the presence of winds, there are still two major open questions: a) the presence of the winds driven by large-scale magnetic field; and b) the formation of a large-scale field. This model predicts strong winds driven by the field, which seem to be absent in most X-ray binaries. But in order to make the winds visible, observationally, they would have to interact with ambient gas or produce internal shocks. So if there aren't internal shocks in the winds, the winds will be invisible due to the lack of ambient gas surrounding the disk in X-ray binaries. But X-ray absorption lines resulting from disk winds may be detected when the mass loss rate is important ($B_{\rm \phi}\sim B_{\rm p}$). Such absorption lines do seem to be present in some X-ray binaries \citep{k2012,m2012b}. The formation of a large-scale field threading a thin disk is a key point in this work. We investigate the advection and diffusion timescales of the field in the disk in section \ref{fields}. Our results are basically the same as those of \citet{c2013}, i.e., the advection timescale can be smaller than the diffusion timescale and the field can be effectively dragged inwards, if the field is initially strong enough. But the main reason for this is that the wind takes away lots of the gravitational energy and so the diffusion timescale becomes larger (Figs. \ref{f6}$-$\ref{f9}). However, even if the field can be effectively dragged inwards, the formation of the large-scale field still depends on the outer boundary conditions. An original large-scale field is needed on the outer boundary of the disk. MHD simulations also suggest the formation of large-scale field depending on the outer boundary conditions \citep{b2008,m2012}. But where the original field comes from is still an unsolved problem. | 14 | 3 | 1403.2305 |
1403 | 1403.1901_arXiv.txt | We present a spectral analysis of three simultaneous \nustar and \swift/XRT observations of the transient Be-neutron star binary \ks taken during its outburst in 2013/2014. These broad-band observations were supported by \swift/XRT monitoring snap-shots every 3\,days, which we use to study the evolution of the spectrum over the outburst. We find strong changes of the power-law photon index, which shows a weak trend of softening with increasing X-ray flux. The neutron star shows very strong pulsations with a period of $P\approx18.8$\,s. The 0.8--79\,keV broad-band spectrum can be described by a power-law with an exponential cutoff and a black-body component at low energies. During the second observation we detect a cyclotron resonant scattering feature at 12.5\,keV, which is absent in the phase-averaged spectra of observations 1 and 3. Pulse phase-resolved spectroscopy reveals that the strength of the feature changes strongly with pulse phase and is most prominent during the broad minimum of the pulse profile. At the same phases the line also becomes visible in the first and third observation at the same energy. This discovery implies that \ks has a magnetic field strength of $B\approx1.1\times10^{12}(1+z)$\,G, which is at the lower end of known cyclotron line sources. | \ks was independently discovered with \textsl{Mir-Kvant}/TTM by \citet{borozdin90a} and with \textsl{CGRO}/BATSE by \citet{finger94a} and \citet{chakrabarty95a} during outbursts in 1989 and 1994, respectively. \citet{swank00a} used \xte data during an outburst in 2000 and the 18.7\,s pulse period to identify both detections as the same accreting neutron star. The optical companion was identified by \citet{negueruela03a} as a Be-type star at a distance of $\sim$10\,kpc, assuming a standard luminosity. \citet{galloway04a} determined the orbit and found an orbital period of $P_\text{orb}=41.5$\,d, with a very low eccentricity of $e=0.034\pm0.007$. In 2000 \xte performed an extensive campaign to monitor a large outburst that reached a peak flux of 120\,mCrab in the 1.5-12\,keV band. \citet{galloway04a} found that the energy spectrum could be described with a simple Comptonization model \citep[\texttt{compTT},][]{titarchuk94a, hua95a}, a model often applied to highly magnetized neutron stars. They found no source-intrinsic absorption, but a broad excess around 10\,keV which they described with a hot black-body component with $kT_\text{bb}=3\text{--}4$\,keV. Using \sax data taken during the decay of the same major outburst, \citet{naik06a} found a similar spectral shape but a much cooler black-body component, $kT_\text{bb}\approx0.6$\,keV. They additionally found evidence for a \feka line at $\sim$6.6\,keV. The major outburst was followed by a series of weaker outbursts, the strongest of which occurred in 2004 April and reached $\sim$45\,mCrab in the 1.5--12\,keV energy band. This series of outbursts was serendipitously monitored by \inte during its Galactic Plane scans. \citet{tsygankov05a} described the \inte/ISGRI and JEM-X spectra using a power-law with a high-energy cut-off and found indications for a spectral softening with increased flux. Accreting neutron-stars sometime show cyclotron resonant scattering features (CRSFs) in their hard X-ray spectra. These absorption-like lines are the only way to directly measure the magnetic field strength close to the neutron star surface. They are produced by photons that scatter off electrons quantized onto Landau-levels in the strong magnetic field ($B\approx10^{12}$\,G) of the neutron star. Their energy is directly related to the strength of the magnetic field in the line forming region via the ``12-B-12''-rule: \begin{equation} E_\text{CRSF}=11.57\times{B}_\text{12}(1+z)\,\text{keV} \end{equation} where $B_{12}$ is the magnetic field in $10^{12}$\,G and $z$ the gravitational redshift \citep[for a detailed discussion see, e.g.,][]{schoenherr07a}. Theoretically CRSF could also result in emission features \citep{schoenherr07a}, but there is only little observational evidence to date \citep[a possible detection was reported for 4U~1626$-$67, see][]{iwakiri12a}. Despite coverage with \xte, \sax, and \inte, a CRSF was not detected in previous outbursts of \ks \citep{naik06a, galloway04a, tsygankov05a}. \begin{deluxetable}{lcccc} \tablecolumns{5} \tablewidth{0pc} \tablecaption{Observation log for the three simultaneous observations.\label{tab:obslog}} \tablehead{\colhead{Observatory/} &\colhead{ ObsID} &\colhead{start date} &\colhead{exposure} & \colhead{pulse period} \\ \colhead {Instrument} & \colhead{} & \colhead{MJD (d)} & \colhead {(ks)} & \colhead{(s)}} \startdata \nustar & 80002015002 & 56586.79 & 18.4 & $18.80584(16)$ \\ \swift/XRT & 00032990003 & 56587.25 & 0.37 & \\ \nustar & 80002015004 & 56618.91 & 18.6 & $18.78399(7)$ \\ \swift/XRT & 00032990013 & 56618.94 & 0.93 & \\ \nustar & 80002015006 & 56635.75 & 25.4 & $18.77088(6)$ \\ \swift/XRT & 00032990020 & 56635.67 & 0.91 & \enddata \end{deluxetable} \ks has been in quiescence from 2004--2013. In 2013 October MAXI \citep{maxiref} detected increased flux levels \citep{Atel5438}. The beginning of an outburst was immediately confirmed by \swift/XRT \citep{Atel5441} and monitored by \swift/BAT. We triggered \swift/XRT $\sim$1\,ks snap-shot observations every 3\,days to monitor the outburst in soft X-rays (Figure~\ref{fig:batlc}). It reached a peak flux of $\sim$130\,mCrab in the 3--10\,keV energy band, very comparable to the maximum of the bright 2000 outburst \citep{naik06a}. Additionally, we triggered three observations with the \textsl{Nuclear Spectroscopy Telescope Array} \citep[\nustar;][]{harrison13a}. An overview of the observations and their exposure times can be found in Table~\ref{tab:obslog}. \begin{figure} \centering \includegraphics[width=0.99\columnwidth]{batlcpar6_bb_4paper_jk.eps} \caption{\textit{(a)} Lightcurve of the 2013 outburst. Data from \swift/BAT \citep[15--50\,keV,][]{swiftbatref} are shown as green crosses, MAXI (2--20\,keV) as blue circles, \swift/XRT (0.5--8\,keV) as black squares and \nustar/FPMA (3--79\,keV) as red diamonds. All data are scaled to the Crab count-rate in the respective energy band. The \nustar count-rate for the \nustar data is shown on the right-hand $y$-axis. \textit{(b)} Best-fit power-law index $\Gamma$ of the \swift/XRT spectra, for details of the model see text. The orange diamonds show combined XRT and \nustar results. } \label{fig:batlc} \end{figure} | \label{sec:outlook} We have presented a spectral analysis of three \nustar observations of the Be-X-ray binary \ks with simultaneous \swift/XRT data, taken during its large 2013/2014 outburst. The broad spectral coverage provided by the combination of these two instruments allowed us to discover a CRSF absorption feature around 12.5\,keV. The feature was significantly detected in the phase-averaged spectrum of the brightest observation, and during the broad pulse minimum in phase-resolved spectroscopy in all observations. During the pulse maximum the feature is not seen significantly, either in absorption or emission. The line energy and width is similar to the lines detected in 4U\,0115+63 and Swift~J1626.6$-$5156 \citep{white83a, decesar13a}. We deduce a surface magnetic field of $\sim1.1\times10^{12}(1+z)$\,G, assuming that the line is the fundamental line. Here $z$ is the gravitational redshift, defined by \begin{equation} (1+z)^{-1}=\sqrt{1-\frac{2GM}{Rc^2}}. \end{equation} For typical neutron star parameters, $z\approx0.3$ if the line-forming region is close to the surface. This magnetic field strength puts \ks at the lower end of known cyclotron lines sources \citep[cf.][]{caballero07a}. During the broad minimum phase of the pulse profile, we detect the CRSF in all three observations. The luminosity near $10^{38}$\,erg\,s$^{-1}$ puts \ks clearly in the super-critical accretion regime, where the radiation pressure is strong enough to decelerate the in-falling matter before the neutron star surface via a radiation-dominated shock \citep{becker12a}. In this regime, a negative correlation between the CRSF energy and luminosity is expected \citep{becker12a}, as observed, for example, in V\,0332+53 \citep{tsygankov10a}. If the correlation were of a similar strength as observed in V\,0332+53 we would not have detected it due to the very small range of luminosities sampled. The time-resolved \swift/XRT spectra show a strongly variable photon-index $\Gamma$ over the outburst, with changes of 10\% or more within 3\,days and softening with increasing X-ray flux. This softening agrees with the expected behavior in the supercritical accretion regime, as shown by \citet{klochkov11a} for various other sources. However, because we restricted the model to describe basically all changes in spectral hardness in the photon-index, it is probable that the true physical changes are more complex than a variable photon-index, e.g., the black-body temperature might vary independently of the X-ray flux. Nonetheless, intrinsic source variability must be present. We clearly detect a \feka line in all data sets, with an energy significantly above the line energy for neutral iron (see Table~\ref{tab:bestfit_all}) and broadened in excess of the energy resolution of \nustar. While Doppler-broadening could be responsible for part of the observed width, the increased energy indicates that the fluorescence region is slightly ionized and the observed broadening originates from a blend of \feka at different low ionization states. The data do not allow us to disentangle different lines from one single broad line. | 14 | 3 | 1403.1901 |
1403 | 1403.8011_arXiv.txt | The quark-meson model is investigated for the two- and three-flavor case extended by contributions of vector mesons under conditions encountered in core-collapse supernova matter. Typical temperature ranges, densities and electron fractions, as found in core-collapse supernova simulations, are studied by implementing charge neutrality and local $\beta$-equilibrium with respect to weak interactions. Within this framework, we analyze the resulting phase diagram and equation of state (EoS) and investigate the impact of undetermined parameters of the model. The EoS turns out to be relatively independent on the entropy per baryon but there are significant changes when going from the two-flavor to the three-flavor case due to the nontrivial contribution from the strange quarks which stay massive even at high densities. While an increasing vector meson coupling constant leads to a substantial stiffening of the EoS, we find that the impact of changing the scalar meson mass is equally strong and results in a softening of the EoS for increasing values. | The properties of strong interaction matter as described by quantum chromodynamics (QCD) at high densities and temperatures can be studied in the laboratory by relativistic heavy-ion collision experiments. Manifestations of this extreme state of matter created in the laboratory can be found in the early universe, neutron star mergers and core-collapse supernova explosions. The processes which are able to turn the collapse of a massive stars into a supernova explosion are not fully understood yet, see e.g.\ \cite{Janka:2012wk} for a review. A key ingredient to core-collapse supernova simulations is the nuclear equation of state at nonzero density, temperature and proton fraction. During the supernova evolution high temperatures and densities can be reached allowing for the opportunity to explore unknown regimes of the phase diagram of strong interactions, i.e.\ the QCD phase diagram. The conditions might be such that a new phase emerges in the core of the collapsed star. A possible phase transition during the supernova evolution has been studied in \cite{Kampfer:1983zz} for a pion-condensed state and in \cite{Takahara:1985,Takahara:1987zq} for a first-order phase transition from hadronic matter to quark matter which can influence the supernova dynamics such that a delayed explosion can take place. Quark matter could also appear during the later proto-neutron star evolution as studied in \cite{Pons:2001ar}. The presence of a new quark matter phase during the supernova stage has been studied in more detail in \cite{Sagert:2008ka,Fischer:2010zz} by including effects from neutrinos. If a new phase is present early in the evolution of the supernova, it can produce a second shock wave with an accompanying measurable second neutrino burst \cite{Dasgupta:2009yj}. In certain cases, different paths in the phase diagram of QCD can be sweeped out by the delayed collapse to a black hole during the evolution of a core-collapse supernova \cite{Ohnishi:2011jv}. The adopted equations of state (EoS) used above are hybrid models with a low-density nucleonic equation of state and the simple MIT bag model extended to nonzero temperatures as the high-density part. The merger of pure quark stars (strange stars) was also simulated within the MIT bag model in \cite{Bauswein:2008gx} showing distinctly different features compared to ordinary neutron star mergers \cite{Bauswein:2009im}. However, it is known that the MIT bag model fails in describing lattice data, see e.g.\ \cite{Fraga:2013qra}, and is not suited to describe profound features of dense matter QCD, as chiral symmetry restoration at high densities. As perturbation theory breaks down on the scale of interest here and results from lattice QCD at high densities are not available yet, improved effective models have to be utilized to study the regime of the QCD plasma relevant for astrophysical applications, as core-collapse supernovae, which we focus on in the following. The Nambu-Jona-Lasinio (NJL) model, as a chiral effective model of QCD, has been studied for nonzero temperature and neutrino chemical potential, relevant for proto-neutron stars, in \cite{Steiner:2002gx,Ruster:2005ib,Laporta:2005be,Sandin:2007zr}. First exploratory investigations of supernova explosions, which require a given electron to baryon number ratio $Y_e$, were only undertaken recently within chiral approaches of QCD in \cite{Fischer:2011zj}. Here the Polyakov-loop extended version of the NJL model, the PNJL model, was used and compared to the MIT bag model. It turned out that there are generic differences between the two model descriptions of relevance for the supernova dynamics. In this work the linear sigma model \cite{GellMann:1960np} is adopted as an effective chiral model of QCD to study supernova matter, where the fundamental particles are quarks interacting via scalar and vector meson exchange. The quark masses are generated by nonvanishing vacuum expectation values of the scalar fields which act as chiral condensates and model the spontaneous breaking of chiral symmetry as observed in the vacuum of QCD. At high temperature and/or densities the condensates melt away and the theory becomes approximately chirally invariant. The order of the phase transition depends on the choice of the parameters used. A detailed study is performed by varying the different parameters of the model to delineate their role for the properties of supernova material. A comparison between the two flavor model, involving only the light up and down quarks, and the three flavor model, where also the strange quark is taken into account, is performed. By using standard methods of finite temperature field theory, the equation of state (EoS) is calculated, which can serve as an input for core-collapse supernova and neutron star merger simulations. | In this work the structure of the phase diagram and the EoS calculated from the linear sigma model in a mean field approximation was analyzed. The model was expanded by adding vector mesons, which give an additional contribution to the quark chemical potentials. Calculations for the two and three flavor cases were presented. The parameters were fixed by measured meson masses and decay constants. The remaining free parameter, like the mass of the $\sigma$ meson $m_{\sigma}$ and the vector coupling constant $g_{\omega}$, were varied and their influences on the phase diagram and the equation of state were investigated. The conditions characteristic for core-collapse supernova explosions served as further input. These are charge neutrality, $\beta-$equilibrium with respect to weak interactions and a given electron-baryon fraction. The phase diagram was analyzed and a first order phase transition at low temperatures was observed for certain parameters. For higher temperatures the phase transition line ends in a critical endpoint, after which the phase transition is a crossover. By increasing the vector meson coupling constant the phase transition gets shifted to higher $\mu _Q$, but for too large values the first order phase transition vanishes and a crossover is observed for all temperatures. The same behavior was seen for increasing $\sigma$ masses. The effect of a lowered $\sigma$ mass could be compensated by an increased value of $g_{\omega}$. The equation of state was calculated for a given entropy per baryon, with typical values found in supernova simulations. The particular value for the entropy per baryon has little influences on the EoS. An increasing vector repulsion, i.e.\ increasing $g_{\omega}$, leads to a higher slope of the EoS. Higher values of $m_{\sigma}$ shift the EoS to lower pressures at constant energy densities. However, the slope of the EoS at high pressures stays constant, whereas at low energy densities differences occur due to the change of the order of the phase transition. Whether the computed equation of state is useful in a supernova explosion has to be tested in simulations. We stress, that the model presented here lacks a low-density hadronic phase. Thus it should be matched at low energy densities to a hadronic EoS for most cases, unless strange quark matter is absolutely stable at vanishing temperature. Recent improvements of the Polyakov-loop extension of the quark-meson model \cite{Haas:2013qwp,Stiele:2013pma,Herbst:2013ufa} can be considered in this framework and the conditions of supernova matter studied here can be worked into the investigation of the nucleation timescales of a quark phase \cite{Mintz:2009ay,Mintz:2012mz}. More importantly, the quark meson model has to be confronted with the observed new mass limit for compact stars from the mass measurement of the pulsars PSR J1614-2230 with a mass of $M=1.97\pm0.04\,\textnormal{M}_{\odot}$ \cite{Demorest:2010bx} and PSR J0348+0432 \cite{Antoniadis:2013pzd} with a mass of $M=2.01\pm0.04\,\textnormal{M}_{\odot}$. This is work in progress and will allow to constrain the parameter space more strictly than it was possible in this analysis \footnote{A. Zacchi, R. Stiele, J. Schaffner-Bielich, in preparation (2014)}. | 14 | 3 | 1403.8011 |
1403 | 1403.6656_arXiv.txt | The Unified Astronomy Thesaurus (UAT) is an open, interoperable and community-supported thesaurus which unifies the existing divergent and isolated Astronomy \& Astrophysics vocabularies into a single high-quality, freely-available open thesaurus formalizing astronomical concepts and their inter-relationships. The UAT builds upon the existing IAU Thesaurus with major contributions from the astronomy portions of the thesauri developed by the Institute of Physics Publishing, the American Institute of Physics, and SPIE. We describe the effort behind the creation of the UAT and the process through which we plan to maintain the document updated through broad community participation. | In astronomy, there have been different initiatives aimed at creating classification systems to be used in the literature. In 1992, Shobbrook and Shobbrook published a 2,551-term thesaurus which was endorsed by the IAU \citep{1992PASAu..10..134S} and named the IAU Thesaurus. Meanwhile, editors from the main astronomy journals developed a parallel system of keywords to characterize and indexed published articles called ``Astronomy Subject Headings'' (or more commonly simply ``journal keywords'') which was adopted by the journals also in 1992. This system, consisting of just over 300 concepts organized in a hierarchy, was considered simpler to use than the IAU Thesaurus and appropriate for creating the annual subject heading index used for browsing the content of the journals. With the passing of time, limitations of this system have become apparent, yet it continues to be used today in all the major astronomy publications. In 2007 the International Virtual Observatory Alliance (IVOA) Semantics Working Group published a study on the use of Vocabularies in the Virtual Observatory (\cite{2009ASPC..411..179G} and \cite{2011arXiv1110.0520D}), providing versions of the IAU Thesaurus and the astronomy subject headings in SKOS (Simple Knowledge Organization System, \citet{skosref}) format. Outside of astronomy, the American Institute of Physics (AIP), on behalf of the publishers in the field, developed a more comprehensive classification system which included astronomy terms. Named PACS (Physics and Astronomy Classification System), the keywords were originally proposed in 1975 and have been used to characterize content in most of the major physics journals until 2011, when AIP announced that they would stop maintaining and using this classification scheme in favor of a more modern system. With the end of PACS, the physics community found itself without a common system for classifying new journal articles, and in the fall of 2011 a group of physics and astronomy publishers met to explore the possibility of joining forces in developing a modern system to support the classification and semantic enrichment of the literature (the new thesaurus's lineage is illustrated in Figure 1). Participants included members of the astronomical community involved in the publications and curation journal articles, including representatives from ADS, the CfA Wolbach Library, and the Institute of Physics (IOP), the publisher of AAS journals. Recognizing the danger of having different publishers develop separate thesauri based on the content that they managed, a key group of participants at the meeting felt that there was an opportunity to collaborate on the development of a single system covering at least astronomy and astrophysics. IOP and AIP, in collaboration with ADS, the CfA Wolbach Library and the IVOA Semantics working group began discussing the possibility of working together on a single astronomy thesaurus created by merging and reconciling existing and emerging vocabularies and thesauri. As the plans to create a unified thesaurus were taking shape, the American Astronomical Society (AAS) joined the effort and provided logistical and legal support helping negotiate the licensing terms for the final work. Having settled issues related to intellectual rights in the fall of 2012, AIP and IOP proceeded to donate the astronomy portions of their thesauri and funded an effort to merge them with the IVOA-enhanced version of the IAU thesaurus. The resulting work, newly named the Unified Astronomy Thesaurus, was born. \begin{figure}[!ht] \plotone{P60_f1.png} \caption{Lineage of the Unified Astronomy Thesaurus: the diagram illustrates the relative size and overlap of the concepts which contributed to the list of terms in the UAT.} \end{figure} | 14 | 3 | 1403.6656 |
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1403 | 1403.4186_arXiv.txt | The BICEP2 experiment determined the Hubble parameter during inflation to be about $10^{14}$\,GeV. Such high inflation scale is in tension with the QCD axion dark matter if the Peccei-Quinn (PQ) symmetry remains broken during and after inflation, because too large axion isocurvature perturbations would be generated. The axion isocurvature perturbations can be suppressed if the axion acquires a sufficiently heavy mass during inflation. We show that this is realized if the PQ symmetry is explicitly broken down to a discrete symmetry and if the breaking is enhanced during inflation. We also show that, even when the PQ symmetry becomes spontaneously broken after inflation, such a temporarily enhanced PQ symmetry breaking relaxes the constraint on the axion decay constant. | The identity of dark matter is one of the central issues in cosmology and particle physics. Among various candidates for dark matter, the QCD axion is a plausible and interesting candidate. The axion, $a$, arises as a pseudo-Nambu-Goldstone (pNG) boson in association with the spontaneous breakdown of a global U(1)$_{\rm PQ}$ Peccei-Quinn (PQ) symmetry~\cite{Peccei:1977hh,QCD-axion}. If the U(1)$_{\rm PQ}$ symmetry is explicitly broken only by the QCD anomaly, the axion is stabilized at vacuum with a vanishing CP phase, solving the strong CP problem. More important, the dynamical relaxation necessarily induces coherent oscillations of axions, which contribute to cold dark matter (CDM). We focus on the axion CDM which accounts for the total dark matter density, throughout this letter. The axion mass receives contributions from the QCD anomaly, \bea \label{maQCD} m^{\rm QCD}_a \simeq 6\times 10^{-6}{\rm eV} \left(\frac{f_a}{10^{12}{\rm GeV}}\right)^{-1}, \eea where $f_a$ is the axion decay constant. Because of the light mass, the axion generically acquires quantum fluctuations of $\delta a \simeq H_{\rm inf}/2\pi$ during inflation, leading to CDM isocurvature perturbations. Here $H_{\rm inf}$ is the Hubble parameter during inflation. The mixture of the isocurvature perturbations is tightly constrained by the Planck observations~\cite{Ade:2013uln}, which reads \bea H_{\rm inf} < 0.87\times 10^7\,{\rm GeV} \left(\frac{f_a}{10^{11}{\rm GeV}}\right)^{0.408}~~~(95\%\,{\rm CL}), \label{iso} \eea neglecting anharmonic effects~\cite{Turner:1985si,Lyth:1991ub,Bae:2008ue,Visinelli:2009zm,Kobayashi:2013nva}. In particular, a large-field inflation such as chaotic inflation~\cite{Linde:1983gd} is in conflict with the isocurvature bound. Recently the BICEP2 experiment announced the discovery of the primordial B-mode polarization~\cite{BICEP2}, which determines the inflation scale as \bea \label{B} H_{\rm inf} &\simeq& 1.0 \times \GEV{14} \lrfp{r}{0.16}{\frac{1}{2}},\\ r &=& 0.20^{+0.07}_{-0.05} ~~(68\%{\rm CL}), \eea where $r$ denotes the tensor-to-scalar ratio.\footnote{Such large tensor-to-scalar ratio can be explained in various large field inflation; see e.g.~\cite{Freese:1990ni,Kawasaki:2000yn,Silverstein:2008sg, McAllister:2008hb,Kaloper:2008fb,Takahashi:2010ky,Nakayama:2010kt,Nakayama:2010sk,Harigaya:2012pg,Croon:2013ana, Nakayama:2013jka,Czerny:2014wza,Czerny:2014xja,Nakayama:2014-HCI}. The tension with the Planck result can be relaxed in the presence of small modulations in the inflaton potential~\cite{Kobayashi:2010pz} or hot dark matter/dark radiation. } After subtracting the best available estimate for foreground dust, the allowed range is modified to $r = 0.16^{+0.06}_{-0.05}$. Therefore one can see from \REF{iso} and \REF{B} that there is a clear tension between the inflation scale determined by the BICEP2 and the QCD axion dark matter.\footnote{ The isocurvature perturbation bound similarly applies to the so called axion-like particles, or general pseudo Nambu-Goldstone bosons, which are produced by the initial misalignment mechanism and contribute to dark matter. } There are various known ways to suppress the axion CDM isocurvature perturbations. First, if the PQ symmetry is restored during inflation (or reheating), there is no axion CDM isocurvature perturbations, as the axion appears only when the PQ symmetry is spontaneously broken some time after inflation~\cite{Linde:1990yj,Lyth:1992tx}. In this case topological defects such as axionic cosmic strings and domain walls are generated, and in particular the domain wall number $N_{\rm DW}$ must be unity to avoid the cosmological catastrophe~\cite{Hiramatsu:2012gg}. Second, if the kinetic term coefficient for the phase of the PQ scalar was larger during inflation than at present, the quantum fluctuations, $\delta a$, can be suppressed after inflation. This is possible if the radial component of the PQ scalar takes a larger value during inflation~\cite{Linde:1990yj,Linde:1991km}. The scenario can be implemented easily in a supersymmetric (SUSY) theory, as the saxion potential is relatively flat, lifted by SUSY breaking effects. Interestingly, a similar effect is possible if there is a non-minimal coupling to gravity~\cite{Folkerts:2013tua}. Third, the axion may acquire a heavy mass during inflation so that its quantum fluctuations get suppressed~\cite{Jeong:2013xta}. In Ref.~\cite{Jeong:2013xta} two of the present authors (KSJ and FT) showed that the QCD interactions become strong at an intermediate or high energy scale in the very early Universe, if the Higgs field has a sufficiently large expectation value.\footnote{ The idea of heavy QCD axions during inflation was considered in Refs.~\cite{Dvali:1995ce,Banks:1996ea,Choi:1996fs} to suppress the axion abundance, not the isocurvature perturbations.} In fact, the second solution of Refs.~\cite{Linde:1990yj,Linde:1991km} is only marginally consistent with the BICEP2 result \REF{B}, if the field value of the PQ scalar is below the Planck scale. Also the third solution of Ref.~\cite{Jeong:2013xta} is consistent with the BICEP2 result only in a corner of the parameter space. Therefore, we need another solution to suppress the axion isocurvature perturbations, as long as we assume that the PQ symmetry remains broken during and after inflation. In this letter, we propose a simple mechanism to suppress the axion CDM isocurvature perturbations along the line of the third solution. Instead of making the QCD interactions strong during inflation, we introduce a PQ symmetry breaking operator, which becomes relevant only during inflation. If the axion acquires a sufficiently heavy mass during inflation, the axion CDM isocurvature perturbations practically vanish, evading the isocurvature bound on the inflation scale. After inflation, the explicit PQ breaking term should become sufficiently small so that it does not spoil the axion solution to the strong CP problem. There are various possibilities to realize such temporal enhancement of the axion mass. If, during inflation, the radial component of the PQ scalar, i.e. the saxion, takes a large field value, the PQ symmetry breaking operator is enhanced and the axion becomes heavy. After inflation, the saxion settles down at the low-energy minimum located at a smaller field value where the explicit PQ breaking term is sufficiently small. Alternatively we can consider a case in which the PQ symmetry breaking operators are present only during inflation. This is the case if the inflaton is coupled to the PQ symmetry breaking operators; during inflation, the operators are enhanced due to a large vacuum expectation value (VEV) of the inflaton, whereas they are suppressed if the inflaton is stabilized at much smaller field values after inflation. This can be nicely implemented in a large-field inflation model where the inflaton field value evolves significantly. Later in this letter we also briefly consider the first solution to the tension between the BICEP2 results and the axion CDM isocurvature perturbations, i.e., the PQ symmetry restoration during or after inflation. We will show that, even in this case, the temporarily enhanced PQ symmetry breaking relaxes the bound on the axion decay constant, allowing $f_a \gtrsim \GEV{10}$ and also $N_{\rm DW} \ne 1$. | We have seen that, if the PQ scalar takes super-Planckian values during inflation, there is an allowed region for $f_a$ where the BICEP2 results become consistent with the axion CDM. Let us here briefly discuss one inflation model with the PQ scalar identified with the inflaton. In non-SUSY case, the PQ scalar can be stabilized by the balance between the negative mass and the quartic coupling as in \EQ{Vps}. Then, it is possible to realize the quadratic chaotic inflation model with the PQ scalar, if the kinetic term of the PQ scalar is significantly modified at large field values, based on the running kinetic inflation~\cite{Takahashi:2010ky,Nakayama:2010kt,Nakayama:2010sk}. For instance, we can consider \bea {\cal L} &=&|\partial S|^2 + \xi (\partial |S|^2)^2 - \frac{\lambda}{4} \left(|S|^2-f^2 \right)^2, \eea where $\xi \gg 1$ denotes the coefficient for the running kinetic term, and $\lambda$ is the quartic coupling. The large value of $\xi$ can be understood by imposing a shift symmetry on $|S|^2$: $|S|^2 \to |S|^2 + C$, where $C$ is a real constant. Then the $\xi$-term respects the symmetry, while the other terms explicitly break the shift symmetry.\footnote{ In other words, the ordinary kinetic term and the potential term are relatively suppressed as they explicitly break the shift symmetry. } Note that the shift symmetry is consistent with the PQ symmetry, as it is the radial component of $S$ that transforms under the shift symmetry. Let us denote the radial component of $S$ as $\sigma = |S|$. At large field values $\sigma \gg 1/\sqrt{\xi}$, the canonically normalized field is ${\hat \sigma}\sim \sqrt{\xi} \sigma^2$, and the scalar potential becomes the quadratic one in terms of $\sigma$. The quadratic chaotic inflation can be realized by the PQ scalar. In this case, $\langle {\hat S} \rangle_{\rm inf}$ is of order $15 M_{Pl}$, and one can see from Fig.\ref{fig:relaxed-constraint} that there is a region between $f_a \simeq \GEV{10}$ and $\GEV{13}$ where the isocurvature constraint becomes consistent with the BICEP2 result. Thus, one interesting way to evade the isocurvature bound is to identify the PQ scalar with the inflaton. In this letter, we have proposed a simple mechanism in which the axion acquires a heavy mass during inflation, leading to the suppression of the axion CDM isocurvature perturbations. The point is that the U(1)$_{PQ}$ symmetry is explicitly broken down to its discrete subgroup, $Z_N$. If the PQ-breaking operators are significant during inflation, the axion can acquire a sufficiently heavy mass, suppressing the isocurvature perturbations. There are two ways to accomplish this. One is that the PQ-breaking operators depend on the inflaton field value. If the inflaton takes a large-field value during inflation, the PQ-breaking operator becomes significant, while it becomes much less prominent if the inflaton is stabilized at smaller field values after inflation. This % possibility nicely fits with the large field inflation suggested by the BICEP2 result \REF{B}. The other is that the saxion field value changes significantly during and after inflation. Then, for sufficiently large power of the PQ symmetry breaking operator, it is possible to realize the heavy axion mass during inflation, while keeping the axion solution to the strong CP problem. The lower bound on the power $N$ reads $N > 12$ for $f_a = \GEV{12}$ and the gravitino mass of order $100$\,TeV. We have discussed a concrete PQ scalar stabilization to show that it is possible that the saxion is fixed at a large field value during inflation, while it settles down at the true minimum located at a smaller field value, without restoring the PQ symmetry. Toward a UV completion of our scenario based on the string theory, we may consider non-perturbative effects, \bea W \supset (\phi -\phi_0)^n e^{\pm{\cal A}}, ~~~{\rm or}~~~ K \supset (\phi^{\dag} -\phi_0^{\dag})^m e^{\pm{\cal A}} + {\rm h.c.}, \eea where the exponentials $\propto e^{\pm{\cal A}}$ break a $U(1)_{\rm PQ}$ symmetry down to a discrete one. Here ${\cal A}$ is a (linear combination of) string theoretic axion multiplet, and $\phi$ is the inflaton. If the inflaton $\phi$ develops non-zero expectation value during the inflation (and $\phi=\phi_0$ in the true vacuum), such an axion obtains a large mass, which may suppress the isocurvature perturbations. Similar terms may also help the overshooting problem of moduli during inflation, producing the high potential barrier against decompactification \cite{Kallosh:2004yh}\footnote{ The so-called Kallosh-Linde model can give a large mass to the relevant axion multiplet: $W \supset (\phi -\phi_0)^n (W_0 + Ae^{\pm a{\cal A}}+ Be^{\pm b{\cal A}})$. See also \cite{He:2010uk}. }, even without the coupling to the inflaton in the superpotential \cite{Abe:2005rx}. In nature there may be other kinds of pseudo Nambu-Goldstone bosons such as axion-like particles, and the isocurvature perturbation bound can be similarly applied to them. The tension between the isocurvature perturbations and the high-scale inflation can be solved by our mechanism. That it to say, we can add a symmetry breaking operator which becomes relevant only during inflation. We can do so by either introducing an inflaton field charged under the symmetry or by assuming that the radial component significantly evolves during and after inflation. In particular, our mechanism can relax the isocurvature bound on the $7$\,keV axion dark matter proposed by the present authors~\cite{Higaki:2014zua} to explain the recently found $3.5$\,keV X-ray line~\cite{Bulbul:2014sua,Boyarsky:2014jta}. {\it Note added:} Isocurvature constraints on the QCD axion and axion like particles were considered also in Refs.~\cite{Marsh:2014qoa,Visinelli:2014twa} soon after the BICEP2 announcement. {\it Note added 2:} After submission of this paper we noticed that some of our arguments has an overlap with Ref.~\cite{Dine:2004cq} where it was pointed out that the axion isocurvature perturbations can be suppressed if the PQ symmetry is badly broken during inflation and the alignment of the axion by the PQ symmetry breaking can eliminate domain walls. | 14 | 3 | 1403.4186 |
1403 | 1403.5075_arXiv.txt | { The Cherenkov Telescope Array (CTA) will be the next high-energy gamma-ray observatory. Selection of the sites, one in each hemisphere, is not obvious since several factors have to be taken into account. Among them, and probably the most crucial, are the atmospheric conditions. Since July 2012, the site working group has deployed automatic ground based instrumentation (ATMOSCOPE) on all the candidate sites. Due to the limited time span available from ground based data, long term weather forecast models become necessary tools for site characterization. It is then of prime importance to validate the models by comparing it to the ATMOSCOPE measurements. We will describe the sources of data (ATMOSCOPE, weather forecasting model and satellite data) for the site evaluation and how they will be used and combined. } | CTA will consist of two arrays of Cherenkov telescopes\cite{ach13}, which aim to: (a) increase sensitivity by an order of magnitude with respect to the existing observatories like H.E.S.S., MAGIC and VERITAS, (b) increase the detection area and hence detection rates, (c) improve the angular resolution and hence the morphology reconstruction of extended sources, and (d) provide energy coverage for photons from some tens of GeV to beyond 100 TeV. The observatory will consist of two arrays: a southern hemisphere array, which allows a deep investigation of galactic sources and of the central part of our Galaxy, but also for the observation of extragalactic objects. The northern hemisphere array is primarily devoted to the study of Active Galactic Nuclei and galaxies at cosmological distances. The arrays will also make contributions to the field of particle physics with searches for dark matter, tests of Lorentz invariance and searches of axion-like particles. CTA will be operated as an open, proposal-driven facility. In order to find the best location for the construction of CTA, a comprehensive site search campaign was conducted. The criteria considered for the site ranking are science performance, which depends on the average annual observation time available and the performance of the array per unit time at a given site, as well as costs and risks. Requirements which are most relevant for the site selection are: \begin{itemize} \item Altitude of the sites between 1500 - 3800m above sea level. \item Area available for the deployment must be $>$10km$^2$ for the South and $>$1km$^2$ for the North. \item Ground slope must be less than 8\%. \item The site should have $>70\%$ of moonless night hours completely cloud free. \item During observations the 10-minutes average wind speed is $<36$ km.h$^{-1}$. \item Ambient air temperature during observations will be -15 to +25$^{\circ}$C. \end{itemize} Following the call for site proposals, a total of nine site proposal were received. Five possible locations in the southern hemisphere are evaluated: two locations in Argentina, two in Namibia and one in Chile. The northern proposals include a site in Mexico, two sites in the USA and a site on the island of Tenerife in Spain. The locations of the nine site sites taken into considerations are given in Table \ref{candidate_site}. The Argentinian sites are described in \cite{all13}, the U.S. sites are presented by \cite{ong13}, and for the Spanish site see \cite{pue13}. | CTA will be the next generation of imaging Cherenkov telescopes with one observatory in the Northern hemisphere and one in the Southern hemisphere. The activities to search for sites for CTA started in 2008. Following a call for site proposals, a total of nine proposal documents were received. Here we have described the basic ideas used for the site search. The site characterization depends on three main sources of information. For most sites, a year of ground based data (ATMOSCOPE data) will be gathered. As this is not sufficient to make a comprehensive comparison of the sites, we must use several long term data sources from sites, such as the numerical weather simulations and the satellite data. A major challenge in evaluating the data is to understand the possible biases and uncertainties. In order to overcome these difficulties, a comparison between the ground based data and the long term data sources is being conducted. This allows to find uncertainties and systematic shifts in the data. The data are also verified for consistency at other nearby sites: at local astronomical observatories and weather stations. In most cases, the comparison between the ground based and other sources of data is satisfactory, yet there are some cases when we can only rely on the limited ground based data. There are also several site related studies carried by the CTA consortium. The potential influence of the presence of the E-ELT lasers in Chile is shown in \cite{gau13b}. There is a mirror test facility in San Antonio de los Cobres in Argentina and is described in \cite{med13}. The CTA site decision will also take into account the estimate of the infrastructure cost. The decision process is currently under way and we expect to have a final recommendation at the end of 2013. \vspace*{0.5cm} \footnotesize{{\bf Acknowledgment: }{ We gratefully acknowledge support from the agencies and organizations listed in this page: http://www.cta-observatory.org/?q=node/22. SV acknowledges support through the Helmholtz Alliance for Astroparticle Particle. }} | 14 | 3 | 1403.5075 |
1403 | 1403.0650_arXiv.txt | We explore energy densities of magnetic field and relativistic electrons in the M87 jet. Since the radio core at the jet base is identical to the optically thick surface against synchrotron self absorption (SSA), the observing frequency is identical to the SSA turnover frequency. As a first step, we assume the radio core as a simple uniform sphere geometry. Using the observed angular size of the radio core measured by the Very Long Baseline Array at 43~GHz, we estimate the energy densities of magnetic field ($U_{B}$) and relativistic electrons ($U_{e}$) based on the standard SSA formula. Imposing the condition that the Poynting power and relativistic electron one should be smaller than the total power of the jet, we find that (i) the allowed range of the magnetic field strength ($B_{\rm tot}$) is $1~{\rm G} \le B_{\rm tot} \le 15~{\rm G}$, and that (ii) $ 1\times 10^{-5} \le U_{e}/U_{B} \le 6 \times 10^{2}$ holds. The uncertainty of $U_{e}/U_{B}$ comes from the strong dependence on the angular size of the radio core and the minimum Lorentz factor of non-thermal electrons ($\gamma_{e,\rm min}$) in the core. It is still open that the resultant energetics is consistent with either the magnetohydrodynamic jet or with kinetic power dominated jet even on $\sim 10$ Schwarzschild radii scale. | \label{sec:intro} Formation mechanism of relativistic jets in active galactic nuclei (AGNs) remains as a longstanding unresolved problem in astrophysics. Although the importance of magnetic field energy density ($U_{\rm B}$) and relativistic electron one ($U_{e}$) for resolving the formation mechanism has been emphasized (e.g., Blandford and Rees 1978), it is not observationally clear whether either $U_{\rm B}$ or $U_{e}$ is dominant at the jet base. Relativistic magnetohydrodynamics models for relativistic jets generally assume highly magnetized plasma at the jet base (e.g., Koide et al. 2002; Vlahakis and Konigl 2003; McKinney and Gammie 2004; Krolik et al. 2005; McKinney 2006; Komissarov et al. 2007; Tchekhovskoy et al. 2011; Toma and Takahara 2013; Nakamura and Asada 2013), while an alternative model assumes a pair plasma dominated ``fireball"-like state at the jet base (e.g., Iwamoto and Takahara 2002; Asano and Takahara 2009 and reference therein). Although deviation from equi-partition (i.e., $U_{e}/U_{B}\approx 1$) is essential for investigation of relativistic jet formation, none has succeeded in obtaining a robust estimation of $U_{e}/U_{B}$ at the jet base. M87, a nearby giant radio galaxy located at a distance of $D_{\rm L}=16.7~{\rm Mpc}$ (Jordan et al. 2005), hosts one of the most massive super massive black hole $M_{\bullet}=(3-6)\times 10^{9}~M_{\odot}$ (e.g., Macchetto et al. 1997; Gebhardt and Thomas 2009; Walsh et al. 2013). Because of the largeness of the angular size of its central black hole, M87 is well known as the best source for imaging the deepest part of the jet base (e.g., Junor et al. 1999). Furthermore, M87 has been well studied at wavelengths from radio to Very High Energy (VHE) $\gamma$-ray (Abramowski et al. 2012; Hada et al. 2012 and reference therein) and causality arguments based on VHE $\gamma$-ray outburst in February 2008 indicate that the VHE emission region is less than $\sim 5\delta~{\rm R_{s}}$ where $\delta$ is the relativistic Doppler factor (Acciari et al. 2009). The Very-Long-Baseline-Array (VLBA) beam resolution at 43~GHz typically attains about $0.21\times 0.43~{\rm mas}$ which is equivalent to $5.3\times 10^{16} \times1.1\times 10^{17} ~{\rm cm}$. When $M_{\bullet}=6\times 10^{9}~M_{\odot}$ holds (Gebhardt et al. 2009), then VLBA beam resolution approximately corresponds to $30\times 60~{\rm R_{s}}$. Recent progresses of Very-Long-Baseline-Interferometry (VLBI) observations have revealed the inner jet structure, i.e., frequency and core-size relation, and distance and core-size relation down to close to $\sim 10$ Schwarzschild radii ($R_{\rm s}$) scale (Hada et al. 2011, hereafter H11). Thus, the jet base of M87 is the best laboratory for investigations of $U_{e}/U_{B}$ in the real vicinity of the central engine. Two significant forward steps are recently obtained in M87 observations which motivate the present work. First, Hada et al. (2011) succeeded in directly measuring core-shift phenomenon at the jet base of M87 at 2, 5, 8, 15, 24 and 43 GHz. The radio core position at each frequency has been obtained by the astrometric observation (H11). Since the radio core surface corresponds to the optically-thick surface at each frequency, the synchrotron-self-absorption (SSA) turnover frequency $\nu_{\rm ssa}$ is identical to the observing frequency itself. \footnote{ Difficulties for applying the basic SSA model to real sources has been already recognized by several authors ( Kellermann and Pauliny-Toth 1969; Burbidge et al. 1974; Jones et al. 1974a, 1974b; Blandford and Rees 1978; Marscher 1987) due to insufficiently accurate determination of $\nu_{\rm ssa}$ and $\theta_{\rm obs}$.} Second, we recently measure core sizes in Hada et al. (2013a) (hereafter H13). Hereafter we focus on the radio core at 43~GHz. In H13, we select VLBA data observed after 2009 with sufficiently good qualities (all 10 stations participated and good uv-coverages). To measure the width of the core, a single, full-width-half-maximum (FWHM) Gaussian is fitted for the observed radio core at 43~GHz in the perpendicular direction to the jet axis and we derive the width of the core ($\theta_{\rm FWHM}$). We stress that the core width is free from the uncertainty of viewing angle. Therefore, using $\theta_{\rm FWHM}$ at 43~GHz, we can estimate values of $U_{e}/U_{B}$ in the 43~GHz core of M87 for the first time. In section 2, we derive an explicit form of $U_{e}/U_{B}$ by using the standard formulae of synchrotron absorption processes. As a first step, we simplify a geometry of the radio core as a single uniform sphere although the real geometry is probably more complicated. In section 3, we estimate $U_{e}/U_{\rm B}$ in the M87 jet base by using the VLBA data at 43~GHz obtained in H13. In section 4, we summarize the result and discuss relevant implications. In this work, we define the radio spectral index $\alpha$ as $S_{\nu}\propto \nu^{-\alpha}$ and we assume $M_{\bullet}=6\times 10^{9}~M_{\odot}$. | Based on VLBA observation data at 43~GHz, we explore $U_{e}/U_{B}$ at the base of the M87 jet. We apply the standard theory of synchrotron radiation to the 43~GHz radio core together with the assumption of a simple uniform sphere geometry. We impose the condition that the Poynting and relativistic electron kinetic power should be smaller than the total power of the jet. Obtained values of $B_{\rm tot}$ and $U_{e}/U_{\rm B}$ are summarized in Tables \ref{table:B} and \ref{table:UeUb} and we find the followings; \begin{itemize} \item We obtain the allowed range of magnetic field strength in the 43~GHz core as $1~{\rm G} \le B_{\rm tot} \le 15~{\rm G}$ in the observed radio core at 43~GHz with its diameter $0.11\--0.20~{\rm mas}$ $(15.5\-- 28.2 ~R_{\rm s})$. Our estimate of $B$ is basically close to the previous estimate in the literature (e.g., Neronov and Aharonian 2007), although fewer assumptions have been made in this work. We add to note that even if $\delta$ of the 43~GHz core becomes larger than unity, the field strength only changes according to $B_{\rm tot}\propto \delta$. It is worth to compare these values with independently estimated $B_{\rm tot}$ in previous works more carefully. Abdo et al. (2009) has estimated Poynting power and kinetic power of the jet by the model fitting of the observed broad band spectrum and derive $B_{\rm tot}=0.055~{\rm G}$ with $R=1.4\times 10^{16}~{\rm cm}=0.058~{\rm mas}$, although they do not properly include SSA effect. Acciari et al. (2009) predict field strength $B_{\rm tot}\sim 0.5~{\rm G}$ based on the synchrotron cooling argument. Since smaller values of $B_{\rm tot}$ lead to smaller $\theta_{\rm obs}$, if we assume $\theta_{\rm obs,min}$ by a factor of $\sim 3$ than the true $\theta_{\rm obs}=0.11~{\rm mas}$, the predicted $B_{\rm tot}$ lies between 0.05 and 0.5 gauss which seems to be in a good agreement with previous work. However, for such a small core, electron kinetic power much exceeds the observed jet power. Our result excludes a strong magnetic field such as $B_{\rm tot}\sim 10^{3-4}~{\rm G}$ which is frequently assumed in previous works in order to activate Blandford-Znajek process (Blandford and Znajek 1977; Thorne et al. 1986; Boldt \& Loewenstein 2000). Although M87 has been a prime target for testing relativistic MHD jet simulation studies powered by black-hole spin energy, our result throw out the caveat that the maximum $B_{\rm tot}$, one of the critical parameters in relativistic MHD jets model, $B_{\rm tot}$ should be smaller than $\sim 15$~G for M87. \item We obtain the allowed region of $U_{e}/U_{B}$ in the allowed $\theta_{\rm obs}$ and $\gamma_{e,\rm min}$ plane. The resultant $U_{e}/U_{B}$ contains both the region of $U_{e}/U_{B}>1$ and $U_{e}/U_{B}<1$. It is found that the allowed range is $1\times 10^{-5}\le U_{e}/U_{B}\le 6\times 10^{2}$. The uncertainty of $U_{e}/U_{B}$ is caused by the strong dependence on $\theta_{\rm obs}$ and $\gamma_{e,\rm min}$. Our result gives an important constraint against relativistic MHD models in which they postulate very large $U_{\rm B}/U_{e}$ at a jet-base (e.g., Vlahakis and Konigel 2003; Komissarow et al. 2007, 2009; Tchekhovskoy et al. 2011). To realize sufficiently magnetic dominated jet such as $U_{\rm B}/U_{e}\sim 10^{3-4}$, relatively large $\gamma_{e,\rm min}$ of the order of $\sim 10^{2}$ and a relatively large $\theta_{\rm obs}$ are required. Thus, the obtained $U_{e}/U_{\rm B}$ in this work gives a new constraint on the initial conditions in relativistic MHD models. \end{itemize} Last, we shortly note key future works. \begin{itemize} \item Observationally, it is crucial to obtain resolved images of the radio cores at 43~GHz with space/sub-mm VLBI which would clarify whether there is a sub-structure or not inside $\sim 16$~Rs scale at the M87 jet base. Towards this observational final goal, as a first step, it is important to explore physical relations between the results of the present work and observational data at higher frequencies such as 86~GHz and 230~GHz (e.g., Krichbaum et al. 2005; Krichbaum et al. 2006; Doeleman et al. 2012). Indeed, we conduct a new observation of M87 with VLBA and the Green Bank Telescope at 86~GHz and we will explore this issue using the new data. Space-VLBI program also could play a key role since lower frequency observation can attain higher dynamic range images with a high resolution (e.g., Dodson et al. 2006; Asada et al. 2009; Takahashi and Mineshige 2010; Dodson et al. 2013). If more compact regions inside the 0.11mas region are found by space-VLBI in the future, then $U_{e}/U_{\rm B}$ in the compact regions are larger than the ones shown in the present work. \item Theoretically, we leave following issues as our future work. (1) Constraining plasma composition (i.e, electron/proton ratio) is one of the most important issue in AGN jet physics (Reynolds et al. 1996; Kino et al. 2012) and we will study it in the future. Roughly saying, inclusion of proton powers ($L_{p}$) will simply reduce the upper limit of $B_{\rm tot}$ because $L_{\rm jet} \approx L_{e}+L_{p}\approx L_{\rm poy}$ would hold. (2) On $\sim 10~R_{\rm s}$ scale, general relativistic (GR) effects can be important and they will induce non-spherical geometry. If there is a Kerr black hole at its jet base, for example following GR-related phenomena may happen; (i) magneto-spin effect which aligns a jet-base along black hole spin, and it leads to asymmetric geometry (McKinney et al. 2013). (ii) the accretion disk might be warped by Bardeen and Peterson effect caused by the frame dragging effect (Bardeen and Peterson 1975; Hatchett et al. 1981). Although a recent research by Dexter et al. (2012) suggests that the core emission is not dominated by the disk but the jet component, the disk emission should be taken into account if accretion flow emission is largely blended in the core emission in reality (see also Broderick and Loeb 2009). We should take these GR effects into account when they are indeed effective. (3) Apart form GR effect, pure geometrical effect between jet opening angle and viewing angle which may cause a partial blending of SSA thin part of the jet. It might also cause non-spherical geometry and inclusion of them is also important. \end{itemize} \bigskip \leftline{\bf \large Acknowledgment} \medskip \noindent We acknowledge the anonymous referee for his/her careful review and suggestions for improving the paper. MK thank A. Tchekhovskoy for useful discussions. This work is partially supported by Grant-in-Aid for Scientific Research, KAKENHI 24540240 (MK) and 24340042 (AD) from Japan Society for the Promotion of Science (JSPS). \footnotesize | 14 | 3 | 1403.0650 |
1403 | 1403.6846.txt | Although it is well known that massive galaxies have central black holes, most of them accreting at low Eddington ratios, many important questions still remain open. Among them, are the nature of the ionizing source, the characteristics and frequencies of the broad line region and of the dusty torus. We report observations of 10 early-type galactic nuclei, observed with the IFU/GMOS spectrograph on the Gemini South telescope, analysed with standard techniques for spectral treatment and compared with results obtained with principal component analysis Tomography (Paper I). We performed spectral synthesis of each spaxel of the data cubes and subtracted the stellar component from the original cube, leaving a data cube with emission lines only. The emission lines were decomposed in multi-Gaussian components. We show here that, for eight galaxies previously known to have emission lines, the narrow line region can be decomposed in two components with distinct line widths. In addition to this, broad H$\alpha$ emission was detected in six galaxies. The two galaxies not previously known to have emission lines show weak H$\alpha$+[N II] lines. All 10 galaxies may be classified as low-ionization nuclear emission regions in diagnostic diagrams and seven of them have bona fide active galactic nuclei with luminosities between 10$^{40}$ and 10$^{43}$ erg s$^{-1}$. Eddington ratios are always $<$ 10$^{-3}$. | \label{sec:intro} Nuclear gas emission emerging from early-type galaxies (ETGs) in the local Universe is associated with low-luminosity active galactic nuclei (LLAGN) in 2/3 of the cases \citep{2008ARA&A..46..475H}, where most of these objects are classified as low ionization nuclear emission regions (LINERs; \citealt{1980A&A....87..152H}). In order to distinguish, at optical wavelengths, between LINERs, Seyfert galaxies, starburst regions and transition objects (TOs), the so-called diagnostic diagrams or BPT diagrams \citep{1981PASP...93....5B}, which compare different emission line ratios, are commonly used. Most recent diagnostic diagrams (e.g., \citealt{1997ApJS..112..315H,2003MNRAS.346.1055K,2006MNRAS.372..961K}) compare the [O III]$\lambda$5007/H$\beta$ ratio with the [O I]$\lambda$6300/H$\alpha$, [N II]$\lambda$6581/H$\alpha$ and ([S II]$\lambda$6716 + [S II]$\lambda$6731)/H$\alpha$ line ratios. The reason for such choices is that these ratios are composed of emission lines that are close enough in wavelength, and therefore, are less affected by dust reddening effects \citep{1987ApJS...63..295V}. Photoionization of LINERs by AGNs was proposed by \citet{1983ApJ...264..105F} and \citet{1983ApJ...269L..37H}. However, other mechanisms may be responsible for LINER emission (e.g. shockwaves - \citealt{1980A&A....87..152H}; post-asymptotic giant branch star populations (pAGBs) - \citealt{1994A&A...292...13B}). In the optical, detection of a broad-line region (BLR), specially in the H$\alpha$ line, is a typical feature of AGNs \citep{1997ApJS..112..391H,2008ARA&A..46..475H}. However, the BLR is not detected in several LINERs \citep{1997ApJS..112..391H}. One hypothesis is that the BLR is very weak in these objects and demands an accurate subtraction of the star light from the spectra \citep{1997ApJS..112..391H}. Although there are no reasons to discard the unified models \citep{1993ARA&A..31..473A} in LINERs (see, for instance, \citealt{1999ApJ...525..673B}), BLRs seem to be intrinsically absent, even if their radio or X-ray emissions are typically non-stellar \citep{2008ARA&A..46..475H}. Because of the low bolometric luminosities and Eddington ratios, which are characteristics of LINERs, the formation of BLRs may not occur in some objects \citep{2000ApJ...530L..65N,2006ApJ...648L.101E,2008ARA&A..46..475H,2011A&A...525L...8C}. This is the second of a series of papers whose goal is to detect and characterize the nuclear and circumnuclear (scales of an order of 100 pc) gas emissions in a sample of 10 ETGs. In Ricci et al. (2014, hereafter Paper I), we analysed this sample with the principal component analysis (PCA) Tomography technique only \citep{1997ApJ...475..173H,2009MNRAS.395...64S}. This methodology consists in applying PCA to data cubes. With PCA Tomography, one is able to detect spatial and spectral correlations along data cubes, allowing an efficient and statistically optimized extraction of information from the observed region. In Paper I, we concluded that at least eight galaxies contain an AGN in their central region. In addition, we showed that seven galaxies possess gas discs in their circumnuclear environment. In one object, the circumnuclear gas component seemed to have an ionization cone structure. Stellar discs, also in the circumnuclear regions, were detected in seven galaxies. In order to validate the results presented in Paper I, we will analyse the data cubes of the sample galaxies with techniques that are well established in the literature. In fact, in this work and in Paper III, we will study data cubes with the stellar components properly subtracted. The subtraction was carried out with the stellar population synthesis in each spectrum of the data cubes. With the gas cubes (i.e., data cubes with the gas component only), one is allowed to analyse and characterize the emission lines contained along the observed field of view (FOV) of the galaxies of the sample. In this paper, we intend to detect and characterize the nuclear emission lines of the 10 ETGs, to unveil LLAGNs in these objects. In Paper I, the eigenspectra related to the AGNs were very similar to LINER spectra. However, eigenspectra display correlations between the wavelengths and, thus, are not adequate to measure emission line fluxes and their respective ratios. Hence, in Paper I, it was not possible to use diagnostic diagrams to classify the emissions accurately. Besides, some parameters, such as the colour excess E(B-V) and the luminosity of the lines are essential for a correct characterization of these regions. Nevertheless, our main intention is not to show that results obtained with PCA Tomography may arise from known methodologies but rather to use both techniques as complementary tools. Joint analysis of the results from this work and from Paper I may highlight useful information that would not be possible to extract using only one of the procedures discussed above. Section \ref{dados_PaperII} presents a brief summary of the general characteristics of the data cubes, in addition to a description of the stellar population synthesis. In section \ref{sec:resuts}, we analyse the spectra extracted from the nuclear regions of the galaxies of the sample. Finally, in section \ref{sec:conc}, we discuss our results and present the main conclusions of this work. | \label{sec:conc} Overall, the results presented in this paper are consistent with those from Paper I. We found that all 10 galaxies of the sample have some nuclear activity. Moreover, we detected emission lines in NGC 1399 and NGC 1404, although both H$\alpha$ luminosities are quite low. This may cause a variance of the emission lines of the same order of magnitude as the variance of the noise in the data cubes of both galaxies. This explains the non-detection of their gas emission with PCA Tomography in Paper I. NGC 1380 and NGC 3136 contain the highest E(B-V) values among the galaxies of the sample. In Paper I, we proposed that the correlation between the emission lines and the red region of the continuum in the eigenspectra that revealed the AGNs could be caused by dust extinction. In ESO 208 G-21, IC 5181, NGC 1380, NGC 3136, NGC 4546 and NGC 7097, the Na D absorption lines were also correlated with the red region of the continuum, thus indicating an extinction (dust) associated with the interstellar gas. Probably in NGC 1380 and NGC 3136, a fraction of correlation between the red region of the continuum and the emission lines in the AGN eigenspectrum is also caused by the nebular extinction. Within the uncertainties, all the galaxies of the sample may be classified as LINERs. Nevertheless, other classifications should not be ignored. For instance, the AGN of NGC 4546 is on the limit between LINERs and Seyferts. In fact, its Eddington ratio is the highest among the galaxies of the sample. Besides, its bolometric luminosity is only lower, within the errors, than the $L_{bol}$ of the AGN of IC 1459. Thus, a Seyfert classification is plausible for NGC 4546, according to the diagnostic diagram and to what is predicted by the $L_{bol}$ and $R_{Edd}$ distributions obtained from the Palomar survey (see fig. 9 of \citealt{2008ARA&A..46..475H}). In six galaxies of the sample, we detected broad H$\alpha$ and H$\beta$ emissions, which corroborates the presence of an AGN in these objects. In all cases, FWHMs $>$ 2000 km s$^{-1}$. \citet{1997ApJS..112..391H} found similar FWHM values for this feature in galaxies from the Palomar survey (see their Table 1). Since these features are weak in LINERs, one should be cautious when characterizing these components. Gaussian fits performed in the [N II]+H$\alpha$ region are similar to what \citet{1997ApJS..112..391H} have done. A difficulty highlighted by \citet{1997ApJS..112..391H} is that, since the narrow components do not show exactly a Gaussian profile, broader wings from the NLR may also contaminate the BLR. This may be happening for NGC 2663, whose NLR Gaussian profiles are the broadest among the galaxies of the sample. Besides, the PSF measured with image from the red wing of the broad H$\alpha$ component is larger than the seeing estimated for the observation of this galaxy (see Paper I). Although the seeing of the observation and the PSF of the data cube are not the same thing, one should expect similar FWHMs for both parameters. Another systematical effect which may affect the analysis of the BLR is a bad starlight subtraction. Notwithstanding, the spatial profiles extracted from the images of the red wings of the BLR reveal point-like objects in these six objects. This agrees with an AGN interpretation. Since the broad components have $V_r$ $>$ 280 km s$^{-1}$, the red wings that emerge from the BLR have a huge fraction of the intensity of this feature. In NGC 1399 and NGC 1404, the H$\alpha$ + [N II] emission lines were detected in their nuclear region. A comparison between H$\alpha$ luminosities and X-ray data from both objects suggests that NGC 1399 and NGC 1404 must contain AGNs with very low luminosities, although their $EW(H\alpha)$ measurements indicate that these emissions are too weak to be produced by AGNs. In fact, the detection of a UV point-like source and also strong radio jets in NGC 1399 \citep{2005ApJ...635..305O,2008MNRAS.383..923S} leave no doubt about the existence of an AGN in this galaxy. In NGC 1404, according to \citet{2011ApJ...731...60G}, the combination between X-ray and IR (2MASS) data results in a bolometric luminosity that is high enough for this galaxy to contain an AGN. Below, we summarize the main conclusions of this work. \begin{itemize} \item With well-established methodology of spectra analysis, we show that the main conclusions of Paper I were duly recovered. This shows that PCA Tomography may be used as an useful tool to extract informations from data cubes. \item We show that all the galaxies of the sample have emission lines in their nuclear regions. Among the 10 galaxies, in two (NGC 1399 and NGC 1404) PCA Tomography was not able to detect the emission lines. With spectral synthesis techniques, however, we show that both galaxies have very low intensity emission lines. This indicates that, for very low luminosities AGNs, stellar spectral synthesis may be more effective than PCA Tomography in order to detect emission lines related, for instance, to AGNs. \item In six galaxies of the sample, a broad H$\alpha$ component is detected, which evidences the existence of a BLR in these objects. This feature proves that these six galaxies have an AGN as central objects and a likely photoionization source. Multi-Gaussian decomposition of the H$\alpha$+[N II] set shows that these six galaxies possess a NLR with two kinematic features: a narrower one with 50 $<$ FWHM $<$ 500 s$^{-1}$ and an intermediate one with 400 $<$ FWHM $<$ 1000 km s$^{-1}$. \item Among the four galaxies without a detected BLR, two may contain multiple objects in their centre: NGC 1380 and NGC 3136. These galaxies also revealed a diffuse emission that may be related to H II regions. \item In the other two galaxies, NGC 1399 and NGC 1404, the H$\alpha$+[N II] emission lines have very low intensity. However, X-ray and radio data indicate that they also must contain AGNs. \item Judging from the emission line luminosities, all 10 galaxies have $L_{bol}$ between 10$^{40}$ erg s$^{-1}$ and 10$^{43}$ erg s$^{-1}$, which corresponds to $R_{edd}$ between 10$^{-6}$ and 10$^{-3}$. These Eddington ratios are significantly smaller than the ones found for Seyfert galaxies, but resemble those of LINER-like emissions \citep{2008ARA&A..46..475H}. \end{itemize} | 14 | 3 | 1403.6846 |
1403 | 1403.7300_arXiv.txt | We examine if an efficient particle acceleration takes place by a magnetic-field-aligned electric field near the light cylinder in a rotating neutron star magnetosphere. Constructing the electric current density with the actual motion of collision-less plasmas, we express the rotationally induced, Goldreich-Julian charge density as a function of position. It is demonstrated that the \lq light cylinder gap', which emits very high energy photons via curvature process by virtue of a strong magnetic-field-aligned electric field very close to the light cylinder, will not arise in an actual pulsar magnetosphere. | The Crab pulsar (PSR~J0534+2200), one of the youngest pulsars in our Galaxy, shows pulsed signals in a very wide energy range from radio to $\gamma$-rays (e.g., see Abdo et al. 2010 for the observation of this pulsar with {\it Fermi} LAT between 100~MeV and 20~GeV). In the highest energy range, the Major Atmospheric Gamma-Ray Imaging Cherenkev (MAGIC) telescope has detected pulsed signals at 25~GeV (Aliu et al. 2008), which was confirmed by the LAT observations (Atwood et al. 2009). Further observations with the MAGIC telescope and the Very Energetic Radiation Imaging Telescope Array System (VERITAS) have shown that this component extends up to 400~GeV (Aleksi\'{c} et al. 2012; Aliu et al. 2011). To explain such pulsed fluxes in the very high energy (VHE) region (i.e., above 100~GeV), Bednarek (2012) proposed the \lq light cylinder (LC) gap' model from the following reasons: The rotationally induced, Goldreich-Julian (GJ) charge density is given by (Goldreich-Julian~1969) $ \rho_{\rm GJ} = -\left[ \mbox{\boldmath$\Omega$}\cdot\mbox{\boldmath$B$}/(2\pi c) \right] \left[ 1-(\varpi/\varpi_{\rm LC})^2 \right]^{-1}, $ where $\mbox{\boldmath$\Omega$}$ denotes the rotation vector of the neutron star, $\Omega=\vert\mbox{\boldmath$\Omega$}\vert$ its rotation frequency, $\mbox{\boldmath$B$}$ the magnetic field at each point, $\varpi$ the distance from the rotation axis, $\varpi_{\rm LC}=c/\Omega$ the radius of the LC measured from the rotation axis, and $c$ the speed of light. If the real charge density, $\rho_{\rm r}$, coincides $\rho_{\rm GJ}$ at every position, the magnetic-field-aligned electric field, $E_\parallel$, vanishes in the entire region of the magnetosphere. If $\rho_{\rm r}$ deviates from $\rho_{\rm GJ}$ at some position, on the other hand, the acceleration electric field $E_\parallel$ will arise around that position. In this expression, $\rho_{\rm GJ}$ appears to diverge at the LC, $\varpi \rightarrow \varpi_{\rm LC}$; thus, it was argued if $\rho_{\rm r}$ inevitably deviates from $\rho_{\rm GJ}$ near the LC. By virtue of this diverging behavior of $\rho_{\rm GJ}$, an extremely strong $E_\parallel$, which is about $10^3$ stronger than what arises in the outer-magnetospheric particle accelerator (or the outer gap), was assumed to arise in the LC gap, and the resultant curvature emission was implied to reproduce the pulsed spectrum observed from the Crab pulsar up to 400~GeV. In \S~\ref{sec:GJ}, we demonstrate that the LC gap model is not feasible, examining the actual $\rho_{\rm GJ}$ distribution. Then in \S~\ref{sec:disc}, we briefly mention an appropriate way to compute $\rho_{\rm GJ}$. \section[]{Goldreich Julian charge density in the outer magnetosphere} \label{sec:GJ} In the special relativistic limit, the GJ charge density is given by (e.g., Mestel \& Wang~1982) \begin{equation} \rho_{\rm GJ} \equiv -\frac{\mbox{\boldmath$\Omega$}\cdot\mbox{\boldmath$B$}} {2\pi c} +\frac{(\mbox{\boldmath$\Omega$}\times\mbox{\boldmath$r$})\cdot (\nabla\times\mbox{\boldmath$B$})} {4\pi c}. \label{eq:def_rhoGJ_1} \end{equation} From the inhomogeneous part of the Maxwell equations, we obtain, \begin{equation} \nabla\times\mbox{\boldmath$B$} = \frac{4\pi}{c} \mbox{\boldmath$J$} +\frac{1}{c}\frac{\partial\mbox{\boldmath$E$}}{\partial t}. \label{eq:divB} \end{equation} Since the plasmas are highly collision-less in a pulsar magnetosphere, charged particles gyrate many times between collisions. Thus, we must construct the electric current $\mbox{\boldmath$J$}$ from the gyrating and drifting motion of charged particles, not from the generalized Ohm's law. Let us decompose the current into the parallel and perpendicular components with respect to the local magnetic field line, \begin{equation} \mbox{\boldmath$J$} = \mbox{\boldmath$J$}_\parallel+\mbox{\boldmath$J$}_\perp. \label{eq:J1} \end{equation} First, we consider the parallel current. Since the radiation force balances with the electrostatic acceleration, particles' distribution becomes mono-energetic. Thus, denoting the terminal velocity of out-going particles (e.g., positrons) with $\mbox{\boldmath$v$}_{\parallel +}$, and in-going ones (e.g., electrons) with $\mbox{\boldmath$v$}_{\parallel -}$, we obtain \begin{equation} \mbox{\boldmath$J$}_\parallel = e( n_+ \mbox{\boldmath$v$}_{\parallel +} -n_- \mbox{\boldmath$v$}_{\parallel -}) \label{eq:J2} \end{equation} where $n_+$ (or $n_-$) denotes the number density of out-going (or in-going) particles, and $\mbox{\boldmath$v$}_{\parallel \pm}$ is given by (Hirotani 2011, ApJ 733, L49) \begin{eqnarray} \mbox{\boldmath$v$}_{\parallel \pm} &=& c f_\pm \frac{\mbox{\boldmath$B$}}{B}, \nonumber\\ f_\pm &\equiv& -\frac{\varpi}{\varpi_{\rm LC}}\frac{B^{\hat\varphi}}{B} \pm \sqrt{1-\left(\frac{\varpi}{\varpi_{\rm LC}}\right)^2 \left(\frac{B_{\rm p}}{B}\right)^2}; \label{eq:J3} \end{eqnarray} $e$ denotes the charge on the out-going particle (presumably the positron), $B_{\rm p}^2=B^2-(B^{\hat\varphi}{})^2$, and $B^{\hat\varphi}$ the toroidal component of the magnetic field. In the higher altitudes (e.g., near the LC), it is reasonable to assume $n_+ \gg n_-$ in the gap. In this case, the real charge density is given by $\rho_{\rm r}=e(n_+-n_-) \approx e n_+$. Thus, we obtain \begin{equation} \mbox{\boldmath$J$}_\parallel = \rho_{\rm r} \mbox{\boldmath$v$}_{\parallel +} \label{eq:J4} \end{equation} Even if $n_- \approx n_+$, the additional term that would appear in the right-hand side will not change the entire discussion of this paper; however, we assume $n_- \ll n_+$ to clarify the logic. Second, we consider the perpendicular current. It is given by \begin{eqnarray} \mbox{\boldmath$J$}_\perp &=& c \rho_{\rm r} \frac{\mbox{\boldmath$E$}\times\mbox{\boldmath$B$}}{B^2} + c(P_\perp+P_\parallel) \frac{\mbox{\boldmath$B$}}{B^2} \times \frac{\nabla B}{B} \nonumber\\ && + \frac{c^2 \rho}{B^2}\dot{\mbox{\boldmath$E$}}_\perp -c\nabla \times \left( P_\perp \frac{\mbox{\boldmath$B$}}{B^2} \right), \label{eq:J5} \end{eqnarray} where $P_\parallel$ and $P_\perp$ denote the pressure associated with the longitudinal and perpendicular motion with respect to the magnetic field; $\rho$ (in the second line) denotes the mass density of the plasmas, and $\dot{\mbox{\boldmath$E$}}_\perp$ the temporal derivative of the electric field projected on the perpendicular plane to $\mbox{\boldmath$B$}$. In the right-hand side, the first term represents the current due to the $E \times B$ drift, the second term the sum of the currents due to the magnetic-gradient and the magnetic-curvature drift, the third term (in the second line) the polarization-drift current, and the last term the magnetization current. In a collision-less plasma, the pressure tensor becomes highly anisotropic. In a pulsar magnetosphere, pairs are created inwards (via photon-photon and/or magnetic pair creation in the middle or lower altitudes) with the typical Lorentz factor of a few thousand. Thus, positrons (or electrons) lose most of their perpendicular momentum when they return outwards by a positive (or a negative) $E_\parallel$ in a strong magnetic field. Moreover, their pitch angles decrease due to a subsequent acceleration by $E_\parallel$, resulting in $P_\perp \ll P_\parallel$. Thus, for particles migrating in the outer magnetosphere, we obtain \begin{equation} \mbox{\boldmath$J$}_\perp = c \rho_{\rm r} \frac{\mbox{\boldmath$E$} \times \mbox{\boldmath$B$}} {B^2} +c P_\parallel \frac{\mbox{\boldmath$B$}}{B^2} \times \frac{\nabla B}{B} + \frac{c^2 \rho}{B^2}\dot{\mbox{\boldmath$E$}}_\perp. \label{eq:J7} \end{equation} In a co-rotating magnetosphere, we can put $\rho_{\rm r}=\rho_{\rm GJ}$ and have $c \mbox{\boldmath$E$} \times \mbox{\boldmath$B$} = (\mbox{\boldmath$\Omega$} \times \mbox{\boldmath$r$}) B^2$. Thus, combining equations~(\ref{eq:def_rhoGJ_1}), (\ref{eq:divB}), (\ref{eq:J1}), (\ref{eq:J4}), and (\ref{eq:J7}), we obtain \begin{eqnarray} \lefteqn{ \left[ 1-\left( f_+ \frac{B^{\hat\phi}}{B} +\frac{\gamma m_{\rm e} c^2}{eB} \frac{b^{\hat\varphi}}{L} +\frac{m_{\rm e}c \dot{E}^{\hat\varphi}} {eB^2} \right) \frac{\varpi}{\varpi_{\rm LC}} -\left(\frac{\varpi}{\varpi_{\rm LC}}\right)^2 \right] \rho_{\rm GJ} } \nonumber\\ &=& -\frac{\mbox{\boldmath$\Omega$}\cdot\mbox{\boldmath$B$}} {2\pi c} +\frac{\mbox{\boldmath$e$}_{\hat\varphi}\cdot\dot{E}} {4\pi c} \frac{\varpi}{\varpi_{\rm LC}}, \label{eq:rho_1} \end{eqnarray} where \begin{equation} \frac{b^{\hat\varphi}}{L} \equiv \mbox{\boldmath$e$}_{\hat\varphi} \cdot \left( \frac{\mbox{\boldmath$B$}}{B} \times \frac{\nabla B}{B} \right), \end{equation} $L \sim \varpi_{\rm LC}$, $\vert b^{\hat\varphi} \vert \sim 1$, and $\mbox{\boldmath$e$}_{\hat\varphi}$ denotes the toroidal unit vector. Note that $\gamma m_{\rm e} c^2$ is much small compared to $eBL$ if particles efficiently radiate, and that $\vert \dot{E}^{\hat\varphi} \vert < \Omega B$. We thus finally obtain \begin{equation} \rho_{\rm GJ} = \frac{\displaystyle{-\frac{\mbox{\boldmath$\Omega$}\cdot \mbox{\boldmath$B$}} {2\pi c} +\frac{\mbox{\boldmath$e$}_{\hat\varphi} \cdot\dot{E}} {4\pi c} \frac{\varpi}{\varpi_{\rm LC}} } } {\displaystyle{ 1-f_+ \frac{B^{\hat\phi}}{B} \frac{\varpi}{\varpi_{\rm LC}} -\left(\frac{\varpi} {\varpi_{\rm LC}} \right)^2 } } \label{eq:rho_2} \end{equation} In the numerator, the second term is usually small compared to the first term. In the denominator, we should notice that $f_+$ is positive definite, provided $B^{\hat\phi}<0$. For example, at the LC, we obtain $f_+=2\vert B^{\hat\varphi}\vert/B$. Thus, we find \begin{equation} - f_+ \frac{B^{\hat\phi}}{B} \frac{\varpi}{\varpi_{\rm LC}} > 0. \label{eq:rho_3} \end{equation} It follows that the denominator of equation~(\ref{eq:rho_2}) does not vanish at the LC, provided that the magnetic field is toroidally bent. Moreover, near the LC, we obtain \begin{equation} - f_+ \frac{B^{\hat\phi}}{B} \frac{\varpi}{\varpi_{\rm LC}} \approx 1. \end{equation} Therefore, the GJ charge density is kept around its Newtonian value, $-\mbox{\boldmath$\Omega$}\cdot\mbox{\boldmath$B$}/(2\pi c)$, even near the LC. We can confirm this result by substituting the solution of the vacuum, rotating dipole magnetic field (Cheng et al. 2000) into equation~(\ref{eq:def_rhoGJ_1}). In figure~\ref{fig:rhoGJ}, we plot $\rho_{\rm GJ}/[\Omega B/(2\pi c)]$ as a function of the distance along each magnetic field line. { The magnetic inclination angle $\alpha$ between the magnetic and rotational axes, is assumed to be $60^\circ$ for the solid, dashed, dotted curves, whereas $0^\circ$ for the dot-dot-dot-dashed one. } The solid (or dashed) curves show the results in the trailing (or leading) side of the rotating magnetosphere. The filled circle denotes the position at which the field line crosses the light cylinder. It is confirmed by this explicit calculation that $\rho_{\rm GJ}$ is kept around its Newtonian value even near the LC. { The conclusion is unchanged for small inclination angles. Adopting the vacuum rotating magnetic dipole solution, we obtain $B^{\hat\phi}=0$ if $\alpha=0^\circ$. As a result, equation~(\ref{eq:rho_2}) appears to give a diverging $\rho_{\rm GJ}$ at the LC. Nevertheless, the vacuum solution gives $\nabla\times\mbox{\boldmath$B$}=0$ when $\alpha=0^\circ$. Thus, equation~(\ref{eq:def_rhoGJ_1}) shows that $\rho_{\rm GJ}$ exhibits no singular behavior at the LC. We plot the case of $\alpha=0^\circ$ as the dot-dot-dot-dashed curve in figure~\ref{fig:rhoGJ}, calculating equation~(\ref{eq:def_rhoGJ_1}) from the vacuum, rotating dipole solution. It follows that $\rho_{\rm GJ}$ does not change rapidly at the LC also for an aligned rotator, as expected. } \begin{figure} \includegraphics[width=8.0cm]{f1.eps} \caption{ Distribution of the dimensionless Goldreich-Julian charge density, $\rho_{\rm GJ}/(\Omega B/2\pi c)$, as a function of the distance along the last-open magnetic field line, for discrete values of azimuthal angles, $\varphi_\ast$, measured counter-clockwise around the magnetic axis on the polar-cap surface. Magnetic inclination angle is assumed to be $60^\circ$ { for the solid, dashed, and dotted curves, while $0^\circ$ for the dot-dot-dot-dashed one. } The solid curves represent $\rho_{\rm GJ}/(\Omega B/2\pi c)$ in the trailing side of the rotating magnetosphere { (from the top, $\varphi_\ast= -45^\circ$, $\varphi_\ast= -90^\circ$, $\varphi_\ast= -135^\circ$, and $\varphi_\ast= -180^\circ$) }, while the dashed ones in the leading side { (from the top, $\varphi_\ast= -45^\circ$, $\varphi_\ast= -90^\circ$, and $\varphi_\ast= -135^\circ$) }. The rotational and magnetic axes, as well as the footpoints of the magnetic field lines of $\varphi_\ast=180^\circ$ and $\varphi_\ast=0^\circ$ at the polar-cap surface, reside on the same meridional plane. From the magnetic pole, the direction $\varphi_\ast=180^\circ$ points the rotation axis, while $\varphi_\ast=0^\circ$ the equator. The filled circle denotes the position at which the distance from the rotation axis becomes the light cylinder radius. } \label{fig:rhoGJ} \end{figure} { Recently, } Bednarek (2012) assumed that $\rho_{\rm GJ}$ becomes as large as $\sim 10^3 \Omega B/(2\pi c)$ in a short length $\sim 10^{-3} \varpi_{\rm LC}$ along the magnetic field line in the vicinity of the LC, and considered the curvature radiation that reproduces the pulsed emissions up to 400~GeV from the Crab pulsar. However, since $\rho_{\rm GJ}$ is kept of the order of $\Omega B/(2\pi c)$ even at the LC, this assumption cannot be justified. In another word, the light cylinder gap, which is suggested to produce the pulsed VHE emission from the Crab pulsar, does not appear in any pulsar magnetosphere. | \label{sec:disc} { We arrive at the conclusion that the Goldreich-Julian charge density does not show any singular behavior at the light cylinder. We may note, in passing, that } the Goldreich-Julian charge density should be computed from equation~(\ref{eq:def_rhoGJ_1}) directly, using the given magnetic field distribution in the three-dimensional rotating magnetosphere, instead of replacing $\nabla\times\mbox{\boldmath$B$}$ with the current (i.e., instead of using eq.~[\ref{eq:rho_2}]). We should notice here that we derive equation~(\ref{eq:def_rhoGJ_1}) only from the Maxwell equation, $\nabla\cdot\mbox{\boldmath$E$}=4\pi \rho_{\rm r}$, and the frozen-in condition, assuming stationarity in the co-moving frame, namely $F_{\mu t}+\Omega F_{\mu \varphi} = -\partial_\mu \Psi(r,\theta,\varphi-\Omega t) $, where $F_{\mu\nu}$ represents the field-strength tensor, $\Psi$ the non-corotational potential, and $\mu=t,r,\theta,\varphi$ (Hirotani~2006). That is, equation~(\ref{eq:def_rhoGJ_1}) holds for arbitrary magnetic field, and is derived irrespective of how the current is constructed, or how the plasmas are collisional or collision-less. Let us briefly perform a thought experiment. If the plasma density is large enough, sufficient collisions allow us to use the generalized Ohm's law to describe the current. In this case, $\rho_{\rm GJ}$ does not diverge at the LC, because the $-(\varpi/\varpi_{\rm LC})^2$ term in the coefficient of $\rho_{\rm GJ}$ in equation~(\ref{eq:rho_1}) comes from the $\mbox{\boldmath$E$}\times\mbox{\boldmath$B$}$ drift, which is not included in the Ohm's law. For example, the magnetohydrodynamic approximation, which uses the Ohm's law to close the equations, shows that all the physical quantities are well-behaved at the LC (e.g., Tchekhovskoy et al.~2013). Next, imagine that the plasmas suddenly escape from the magnetosphere to become collision-less. Even in this case, $\rho_{\rm GJ}$ should not be changed at all, because $\rho_{\rm GJ}$ is determined only by the $\mbox{\boldmath$B$}$ field through equation~(\ref{eq:def_rhoGJ_1}), independently from the collisional status of plasmas. Thus, $\rho_{\rm GJ}$ does not diverge at the LC also in the collision-less limit. For example, in the force-free limit, which adopts the $\mbox{\boldmath$E$}\times\mbox{\boldmath$B$}$ drift in $\mbox{\boldmath$J$}_\perp$, no physical quantities diverge or rapidly change at the LC (e.g., Spitkovsky 2006), except for the current sheet, in which the force-free approximation breaks down. In general, quantities behave normally across the LC, without showing any divergence or quick variations. Thus, the light cylinder gap, which has an extreme acceleration electric field (as $10^3$ times stronger than the outer gap), will not arise in a pulsar magnetosphere, unfortunately. | 14 | 3 | 1403.7300 |
1403 | 1403.2713_arXiv.txt | Detection of life on other planets requires identification of biosignatures, i.e., observable planetary properties that robustly indicate the presence of a biosphere. One of the most widely accepted biosignatures for an Earth-like planet is an atmosphere where oxygen is a major constituent. Here we show that lifeless habitable zone terrestrial planets around any star type may develop oxygen-dominated atmospheres as a result of water photolysis, because the cold trap mechanism that protects \ce{H2O} on Earth is ineffective when the atmospheric inventory of non-condensing gases (e.g., \ce{N2}, \ce{Ar}) is low. Hence the spectral features of \ce{O2} and \ce{O3} alone cannot be regarded as robust signs of extraterrestrial life. | The rapid growth of exoplanet discovery and characterization over the last two decades has fueled hopes that in the relatively near future, we may be able to observe the atmospheres of Earth-like planets spectroscopically. Such targets will be intrinsically interesting for comparative planetology, but also for the major reason that they may host life. To search for life on exoplanets by observing their atmospheres, we must first decide on spectral features that can be used as biosignatures. Despite extensive theoretical study of various possibilities, detections of molecular oxygen (\ce{O2}) and its photochemical byproduct, ozone (\ce{O3}), are still generally regarded as important potential indicators of Earth-like life on another planet \citep{Segura2005,Kaltenegger2010,Snellen2013,Kasting2013}. Various authors have investigated the idea that abiotic oxygen production could lead to `false positives' for life \citep{Selsis2002,Segura2007,Leger2011,Hu2013,Tian2014}. For example, it has recently been argued that the build-up of \ce{O2} to levels of $\sim2-3\times10^{-3}$ molar concentration in \ce{CO2}-rich atmospheres could occur for planets around M-class stars, because of the elevated XUV/NUV ratios in these cases \citep{Tian2014}. Extensive atmospheric \ce{O2} buildup due to \ce{H2O} photolysis followed by H escape may also occur on planets that enter a runaway greenhouse state \citep{Ingersoll1969,Kasting1988,Leconte2013b}. However, because by definition the runaway greenhouse only occurs on planets inside the inner edge of the habitable zone, it should not lead to identification of false positives for life. For planets inside the habitable zone, it is commonly believed that \ce{H2O} photolysis will always be strongly limited by cold-trapping of water vapour in the lower atmosphere. The purpose of this note is to point out that a mechanism for \ce{O2} build-up to levels where it is the \emph{dominant} atmospheric gas exists for terrestrial\footnote{Here we define `terrestrial' in the standard (broad) way as describing any planet of low enough mass that it does not possess a dense hydrogen envelope.} planets in the habitable zone around any star type. The reason for this is that the extent of \ce{H2O} cold-trapping depends strongly on the amount of non-condensible gas in the atmosphere. | Because \ce{O2} can become the dominant gas in the atmosphere of a lifeless planet, alone it cannot be regarded as a robust biosignature. Our results do not necessarily rule out its utility in every case. However, they do demonstrate that the situation is considerably more complex than has previously been believed, with the likelihood of an abiotic \ce{O2}-rich atmosphere emerging a complicated function of a planet's accretion history, internal chemistry, atmospheric dynamics and orbital state. Investigation of the range of possibilities for terrestrial planets with variable \ce{N2} and noble gas inventories should be a rich area for future theoretical research that will help to expand our understanding of climate evolution mechanisms. Nonetheless, for a specific exoplanet, even detailed modelling might not lead to a definite conclusion given the inherent uncertainties in processes such as volatile delivery during formation. Observationally, there may still be a way to distinguish the scenarios we discuss here, but only if a reliable way is developed to retrieve the ratio of \ce{O2} to \ce{N2} or \ce{Ar} in an exoplanet's atmosphere. In principle this may be achieved by analysis of the planet's spectrally resolved phase curve \citep{Selsis2011}, or in transit by measurement of the spectral Rayleigh scattering slope \citep{Benneke2012} in a clear-sky (i.e., aerosol-free) atmosphere, or possibly via spectroscopic observation of oxygen dimer features \citep{Misra2014}. More work will be required to assess the potential of these techniques to determine \ce{O2}/\ce{N2} mixing ratios in realistic planetary atmospheres. | 14 | 3 | 1403.2713 |
1403 | 1403.4245_arXiv.txt | Cosmic Reionization On Computers (CROC) is a long-term program of numerical simulations of cosmic reionization. Its goal is to model fully self-consistently (albeit not necessarily from the first principles) all relevant physics, from radiative transfer to gas dynamics and star formation, in simulation volumes of up to 100 comoving Mpc, and with spatial resolution approaching $100\dim{pc}$ in physical units. In this method paper we describe our numerical method, the design of simulations, and the calibration of numerical parameters. Using several sets (ensembles) of simulations in $20h^{-1}\dim{Mpc}$ and $40h^{-1}\dim{Mpc}$ boxes with spatial resolution reaching $125\dim{pc}$ at $z=6$, we are able to match the observed galaxy UV luminosity functions at all redshifts between 6 and 10, as well as obtain reasonable agreement with the observational measurements of the Gunn-Peterson optical depth at $z<6$. | \label{sec:intro} Study of cosmic reionization has been highlighted by the last decadal survey as one of the most promising areas of astrophysical research in the current decade. Progress in this area directly influences many other fields of astrophysics, from thermal evolution of the Lyman-$\alpha$ forest to properties of early galaxies. Because the observational constraints on reionization are limited, theoretical modeling, including numerical simulations, play a relatively larger part in reionization studies than in many other fields of modern astrophysics. Historically, simulations of reionization were mostly confined to two opposite limits: simulations of small spatial volumes with detailed treatment of relevant physics, or large volume simulations with simplified physical modeling \citep[][and \protect\citet{ng:gt09} for a review of the earlier work]{rei:ipm09,rei:at10,rei:fma11,rei:ais12,rei:sim12}. Both approaches suffer from serious limitations. Small box simulations can model individual ionizing sources with sufficient physical detail, but fail to account for the large-scale correlations between them. Large box simulations include these correlations, but, by ignoring gas dynamics, are not able to model ionizing sources self-consistently. The inability of the simulations to include all relevant scales resulted in a recent surge in semi-numerical and purely analytical approximate methods \citep{reisam:fhz04,reisam:fo05,reisam:cf05,reisam:fmh06,reisam:cf06,reisam:zlm07,reisam:mf07,reisam:aa07,reisam:sv08,reisam:zmm11,reisam:mcf11,reisam:vb11,reisam:mfc11,reisam:kf12,reisam:aa12,reisam:mcf12,reisam:zgl13,reisam:btc13,reisam:rfs13,ng:kg13,reisam:sm14}. That's where Moore's Law comes to the rescue. The unrelenting exponential increase in the supercomputing power means that sooner or later the gap between small- and large-box simulations is going to be bridged. In fact, \emph{this time is now} - the new generation of supercomputing platforms that have recently been and are planned to be deployed in the US\footnote{For example, ``Stampede'' at Texas Advanced Computing Center, ``Kraken'' at Oak Ridge National Lab, ``Hopper'' and ``Edison'' at Livermore-Berkeley Lab, ``Mira'' at Argonne National Lab, ``Blue Waters'' at NCSA, etc.}, the so-called ``peta-scale'' platforms (since they get close to or exceed $10^{15}$ floating-point-operations per second), are particularly suitable for large-scale simulations of reionization that treat fully self-consistently the radiative transfer of ionizing radiation and gas dynamics. Taking advantage of this technological progress, we have started a Cosmic Reionization On Computers (CROC) project that aims, over the course of several years, to produce numerical simulations of reionization that model fully self-consistently (albeit not necessarily from the first principles) all relevant physics, from radiative transfer to gas dynamics and star formation, in simulation volumes of up to 100 comoving Mpc and with spatial resolution approaching $100\dim{pc}$ in physical units. In this first paper in a series, we focus primarily on the technical aspects of our simulations, such as the description of the numerical method, simulation design, and the calibration of simulation parameters. We present the original scientific results from our simulations in the subsequent publications. | \label{sec:con} Cosmic Reionization On Computers (CROC) project is a long-term simulation campaign for modeling the process of cosmic reionization in sufficiently large simulations volumes (above 100 Mpc) with detailed physical modeling and spatial resolution better than 0.5 kpc (simulation cell size of less than 150 pc). A simple model of star formation and feedback, based on the linear Kennicutt-Schmidt relation in the molecular gas and a widely used ``delayed cooling'' or ``blastwave'' feedback, is able to reproduce the observed galaxy UV luminosity functions in the whole redshift range $z=6-10$ with a value for the molecular gas depletion time of $\tsf=1.5\dim{Gyr}$, consistent with observations at $z\sim0$ and $z=1-2$. A reasonable choice for the $\euv$ parameter, a quantity that describes photon losses on scales unresolved in our simulations, of $\euv=0.1-0.2$ results in reionization history that is reasonably consistent with the observed opacity of the \lya\ forest in the spectra of SDSS quasars at $z<6$. An even better consistency can be achieved by fitting the simulations to the data, albeit at (presently unrealistically) large computational expense. The observed increase in the Gunn-Peterson optical depth at $z>6$ \citep{igm:fsbw06} has been interpreted by several groups (including ours) as evidence for the reionization overlap \citep{igm:bfws01,igm:wbfs03,ng:g04,igm:fck06,ng:gf06}. That increase is, however, a subject to large cosmic variance; with 6 independent realizations, each corresponding to multiple lines of sight, we find a spread in the redshift of overlap of about $\Delta z\approx1$. Since the observational constraints have even less statistical power than our simulations, they have not yet converged on the true evolution of the average Gunn-Peterson optical depth, and may, therefore, be significantly biased. | 14 | 3 | 1403.4245 |
1403 | 1403.2748_arXiv.txt | Small oscillations of current-carrying string loops around stable equilibrium positions corresponding to minima of the effective potential in the equatorial plane of the Kerr black holes are studied using the perturbation method. In the lowest approximation, two uncoupled harmonic oscillators are obtained that govern the radial and vertical oscillations; the higher-order terms determine non-linear phenomena and transition to chaotic motion through quasi-periodic stages of the oscillatory motion. The radial profiles of frequencies of the radial and vertical harmonic oscillations that are relevant also in the quasi-periodic stages of the oscillatory motion are given, and their properties in dependence on the spin of the black holes and the angular momentum and tension of the string loops are determined. It is shown that the radial profiles differ substantially from those corresponding to the radial and vertical frequencies of the geodetical epicyclic motion; however, they have the same mass-scaling and their magnitude is of the same order. Therefore, we are able to demonstrate that assuming relevance of resonant phenomena of the radial and vertical string-loop oscillations at their frequency ratio $3:2$, the oscillatory frequencies of string loops can be well related to the frequencies of the twin high-frequency quasi-periodic oscillations (HF QPOs) observed in the microquasars GRS 1915+105, XTE 1550-564, GRO 1655-40. We can conclude that oscillating current-carrying string loops have to be considered as one of the possible explanations of the HF QPOs occurring in the field of compact objects. | % Relativistic current-carrying string loops moving axisymmetrically along the symmetry axis of the Kerr or \Schw\nnd\dS{} black holes have been recently studied extensively \cite{Jac-Sot:2009:PHYSR4:,Kol-Stu:2010:PHYSR4:,Stu-Kol:2012:PHYSR4:,Kol-Stu:2013:PHYSR4:}; the current-carrying string loops were first studied in \citep{Lar:1994:CLAQG:,Fro-Lar:1999:CLAQG:}. Tension of such string loops prevents their expansion beyond some radius, while their worldsheet current introduces an angular momentum barrier preventing them from collapsing into the black hole. There is an important possible astrophysical relevance of the current-carrying string loops \cite{Jac-Sot:2009:PHYSR4:} as they could in a simplified way represent plasma that exhibits associated string-like behavior via dynamics of the magnetic field lines in the plasma \cite{Chri-Hin:1999:PhRvD:,Sem-Dya-Pun:2004:Sci:} or due to thin isolated flux tubes of magnetized plasma that could be described by an one-dimensional string \cite{Spr:1981:AA:,Sem-Ber:1990:ASS:,Cre-Stu:2013:PhRvE:}. Motion of electrically charged string loops in combined external gravitational and electromagnetic fields has been recently studied for a Schwarzschild black hole immersed in a homogeneous magnetic field \cite{Arm-etal:2013:PRD:}. The astrophysical applications of the current-carrying string loops have been focused on the problem of acceleration of string loops due to the transmutation process when energy of the oscillatory motion of the string is converted to energy of its translational motion \cite{Jac-Sot:2009:PHYSR4:}. Since the string loops can be accelerated to ultra-relativistic velocities in the deep gravitational potential well of compact objects \cite{Stu-Kol:2012:PHYSR4:,Stu-Kol:2012:JCAP:}, the string loop transmutation can be well considered as a process of formation of ultra-relativistic jets, along with the standard model based on the Blandford-Znajek process \cite{Bla-Zna:1977:MNRAS:} and recently introduced "geodesic collimation" model \cite{Gar-etal:2010:ASTRA:,Pach-etal:2012:ApJ:,Gar-etal:2013:ApJ:}. It has been demonstrated that the cosmic repulsion plays an important role in the acceleration process of the string loops behind the so called static radius of the central object \cite{Stu:1983:BAC:,Stu-Hle:1999:PHYSR4:,Stu-Kol:2012:PHYSR4:}. Here we concentrate our attention on the inverse situation of small oscillations of string loops in vicinity of stable equilibrium positions in the equatorial plane of black-hole spacetimes that was proposed as a possible model of HF QPOs observed in black hole and neutron star binary systems \cite{Stu-Kol:2012:JCAP:}. In the~black hole systems observed in both Galactic and extragalactic sources, strong gravity effects have a~crucial role in three phenomena related to the~accretion disc that is the~emitting source: the~spectral continuum, spectral profiled lines, and oscillations of the~disc; clearly, strong gravity has an~important role also in the~binary systems containing neutron (quark) stars. The best signature of the processes occurring in the strong gravity is frequency of observed oscillations because of possibility to obtain very high precision of its measurements \cite{Fer-etal:2012:ExpAstr:}. HF QPOs of X-ray brightness had been observed in many Galactic Low Mass X-Ray Binaries (LMXB) containing neutron~stars \citep[see, e.g.,][]{Kli:2000:ARASTRA:,Bar-Oli-Mil:2005:MONNR:,Bel-Men-Hom:2007:MONNR:BriNSQPOCor} or black holes \citep[see, e.g.,][]{McCli-Rem:2004:CompactX-Sources:,Rem:2005:ASTRN:,Rem-McCli:2006:ARASTRA:,McCli-Rem:2011:CLAQG:}. Some of the~HF~QPOs are in the~kHz range and often come in pairs of the~upper and lower frequencies ($\nu_{\mathrm{U}}$, $\nu_{\mathrm{L}}$) of {\it twin peaks} in the~Fourier power spectra. The~resonance orbital model of HF QPOs in black hole systems \cite{Tor-etal:2005:ASTRA:}, based on the frequencies of the geodetical orbital and epicyclic motion, is now partially supported by observations, in particular when frequency ratio~3\,:\,2 ($2\nu_{\mathrm{U}} = 3\nu_{\mathrm{L}}$) is seen in twin peak QPOs in the~LMXB containing black holes (microquasars), namely GRO 1655-40, XTE 1550-564, GRS 1915+105 \cite{Tor-etal:2005:ASTRA:}. However, in the case of the GRS 1915+105 source the HF QPO frequency set is more complex - in fact, at least five HF QPOs were observed there \cite{McCli-etal:2006:ARAA:}. Therefore, in this case more complex models of HF QPOs have to be considered. In fact, the complete observed frequency set can be explained in the framework of the extended resonant orbital model (\cite{Stu-Sla-Tor:2007a:ASTRA:,Stu-Sla-Tor:2007b:ASTRA:}) based on the so called Aschenbach effect (\cite{Asch:2004:ASTRA:,Stu-Sla-Tor-Abr:2005:PHYSR4:}); another possibility is related to the multi-resonance orbital model recently proposed in \cite{Stu-Kot-Tor:2013:ASTRA:}. Nevertheless, there remains a clear problem with explanation of the $3:2$ frequency ratios observed in all three microquasars, GRO 1655-40, XTE 1550-564, GRS 1915+105, if the resonance orbital model with a unique variant of the twin oscillating modes with geodetical frequencies is applied, especially when the limits on the black hole spin given by the spectral continuum fitting (\cite{Rem-McCli:2006:ARASTRA:,McCli-Rem:2011:CLAQG:}) are taken into account \cite{Tor-etal:2011:ASTRA:,Ali-etal:2013:CLAQG:,Ste-Gyu-Yaz:2013:PHYSR4:}. It seems that neither the resonant orbital (geodetical) model or any other proposed model could work simultaneously for all of the three microquasars \cite{Tor-etal:2011:ASTRA:}. Therefore, we will test, whether the frequencies of the $3:2$ twin peak oscillations observed in the three microquasars can be explained by the axisymmetric current-carrying string loops oscillating in the field of a Kerr black hole, if the oscillations occur at the "resonant" radii where the radial and vertical frequencies have the rational ratio $3:2$. In Section 2 we summarize the Hamiltonian formalism for the string loop motion and give the perturbative form of the Hamiltonian for motion in vicinity of equilibrium points located in the equatorial plane of Kerr black holes. In Section 3 we give the radial profiles of the frequencies of the radial and vertical modes of the oscillatory motion in terms of the dimensionless spin of the black hole and the angular momentum parameter of the string loop. We discuss their properties, comparing them to the radial profiles of the geodesic radial and vertical epicyclic motion in the Kerr backgrounds. In Section 4 we apply the oscillatory string-loop model to explain the frequencies of the twin peak HF QPOs observed with the frequency ration $3:2$ in the three microquasars. Concluding remarks are presented in Section 5. | We have studied frequencies of harmonic radial and vertical oscillations of relativistic current-carrying string loops in the equatorial plane of Kerr black holes, comparing them to the analogous frequencies of the radial and vertical geodetical epicyclic oscillatory motion of test particles. Magnitudes of both the string loop and test particle frequencies are comparable, having the same mass scaling, however, their radial profiles differ significantly. While the radial epicyclic frequency is always smaller than the vertical epicyclic frequency, the frequencies of the radial and vertical oscillations of the string loops coincide at a "coincidence" radius; the radial frequency overcomes the vertical frequency at radii exceeding the coincidence radius and is smaller under the coincidence radius. The string loop oscillations could be related to the twin HF QPOs observed in microquasars, if we assume their occurrence at resonant radii, where the ratio of the frequencies becomes rational, and a parametric resonance can be relevant. In such a case, we have to consider a dichotomy introduced by the behaviour of the radial profiles of the radial and vertical string loop oscillations, since two resonant points with the same frequency ratio arise for string loop oscillations. In the field of Kerr black holes, the frequencies of the harmonic oscillations of string loops depend on the parameter $\omega$ belonging to the interval $\langle-1,1\rangle$. that reflects the axial angular momentum parameter combining effects of both the angular momentum and tension related to the string, and representing a constant of the string-loop motion. The role of the $\omega$ parameter increases with increasing spin of the black hole, however, for the Schwarzschild spacetime ($a=0$) degenerate situation arises, as the string loop motion is independent of $\omega$. Note that the interplay of the black hole spin $a$ and the string loop parameter $\omega$ has a crucial role in explaining the HF QPOs as the range of allowed frequencies governed by the whole interval of allowed values of $\omega$ increases strongly with increasing spin (see Figs. \ref{psQPO}. and \ref{QPOfit}.) extending thus the ability to cover the observed HF QPO phenomena. In the quasi-periodic regime the frequencies of the string-loop harmonic oscillations are relevant \cite{Stu-Kol:2012:JCAP:,Kol-Stu:2013:PHYSR4:}, therefore, we applied the model of the string-loop oscillations resonant at radii corresponding to the frequency ratio $\nu_{\mit} : \nu_{\mir} = 3:2, 2:3$ to the twin HF QPOs observed with frequency ratio $\nu_{\rm U} : \nu_{\rm L} = 3:2$ in three microquasars GRO 1655-40, XTE 1550-564 and GRS 1915+105. We have demonstrated that the model of the string-loop resonant oscillations can well explain the observed twin HF QPOs for whole the range of the spin parameters expected in the three microquasars and implies a restriction of the range of the mass parameter in microquasars XTE 1550-564 and GRS 1915+105. The observational data put significant restrictions on the string-loop parameter $\omega$ in the case of all three microquasars. The string-loop resonant oscillations model can thus suffice a satisfactory explanation of the twin HF QPOs observed in microquasars, and can give an additional restriction on the black hole parameters. The radial profile of the radial and vertical string-loop oscillations, demonstrating also the case of a unique observed frequency occurring at the coincidence radius, enables potentially explanation of more complex frequency patterns observed in some microquasars, or even in some binary systems containing a neutron star. The test of applicability of the string-loop resonant oscillations model in the cases when more than two HF QPOs are observed, as is the case of the GRS 1915+105 microquasar, or of some other sources, is of high relevance. Such situations will be studied in future works. We conclude that the model of string-loop resonant oscillations can be relevant for the explanation of the HF QPOs observed in microquasars and binary systems containing a neutron (quark) star, and it deserves further investigation related to a more detailed physical description of the properties of the string loops and their oscillations, and models of optical phenomena related to radiating string loops in oscillatory motion. | 14 | 3 | 1403.2748 |
1403 | 1403.7230.txt | We have commenced a detailed analysis of the known sample of Galactic post-asymptotic giant branch (PAGB) objects compiled in the Toru\'n catalogue of Szczerba et al., and present, for the first time, homogeneously derived distance determinations for the 209 \emph{likely} and 87 \emph{possible} catalogued PAGB stars from that compilation. Knowing distances are essential in determining meaningful physical characteristics for these sources and this has been difficult to determine for most objects previously. The distances were determined by modelling their spectral energy distributions (SED) with multiple black-body curves, and integrating under the overall fit to determine the total distance-dependent flux. This method works because the luminosity of these central stars is very nearly constant from the tip of the AGB phase to the beginning of the white-dwarf cooling track. This then enables us to use a standard-candle luminosity to estimate the SED distances. For Galactic thin disk PAGB objects, we use three luminosity bins based on typical observational characteristics, ranging between 3500 and 12000\,$L_{\odot}$. We further adopt a default luminosity of 1700\,$L_{\odot}$ for all halo PAGB objects. We have also applied the above technique to a further sample of 69 related nebulae not in the current edition of the Toru\'n catalogue. In a follow-up paper we will estimate distances to the subset of RV Tauri variables using empirical period-luminosity relations, and to the R\,CrB stars, allowing a population comparison of these objects with the other subclasses of PAGB stars for the first time. | %*********************************************************************** Pre-planetary nebulae (PPNe) are a very brief phase in the late-stage evolution of mid-mass stars ($\sim$1 -- 8\,$M_{\odot}$) between the asymptotic giant branch (AGB) and the planetary nebula (PN) phases (Kwok, Purton \& Fitzgerald 1978; Kwok 1982; Balick \& Frank 2002). The ejection of the tenuous envelope in the final superwind stage of AGB evolution (Renzini 1981) reaches rates of up to $10^{-4}~\msun~\text{yr}^{-1}$, and leads to an increase in effective temperature of the central star. This rate of temperature increase is a strong function of the core mass (Sch\"onberner 1983; Vassiliadis \& Wood 1994, hereafter VW94) and ultimately determines if the core reaches a temperature high enough to photoionize the ejected matter as a planetary nebula (PN), before it disperses into the surrounding interstellar medium (ISM). %%Footnote from SM08 -- The term ÔÔprotoplanetary nebulaeÕÕ (PPNs) is also often used to describe objects in transition between the AGB and PN phases. The relative scarcity of known Galactic PAGB objects ($\sim$450; Szczerba et al. 2007, 2012)\footnote{Earlier compilations of PAGB stars were provided by Szczerba et al. (2001) and Kohoutek (2001).} stems from the brevity of the PAGB evolutionary stage (decades to a few thousand years), which for high core masses can be so brief that we are unlikely to observe these rapidly evolving objects (VW94; Bl\"ocker 1995). The evolution of PPNe is typically characterised by a near constant bolometric luminosity, and a double-peaked spectral energy distribution (SED), manifest as a large infrared excess. Understanding these objects is dependent on accurate distances, which are not available for most of the more poorly-quantified objects. Yet this phase is key to comprehending the shaping mechanisms of PNe (Balick \& Frank 2002), as the dust shells around AGB stars, the precursors of PNe, have morphologies that are typically spherically symmetric (Corradi et al. 2003; Mauron \& Huggins 2006; Cox et al. 2012; Mauron, Huggins \& Cheung 2013), while imaging surveys of PNe show round morphologies to be in the minority (Balick 1987; Manchado et al. 1996; G\'orny et al. 1999; Parker et al. 2006; but see Jacoby et al. 2010). In order to better understand this conundrum, several imaging studies of PPNe have been undertaken over the last two decades, both at optical and infrared wavelengths (Sahai \& Trauger 1998; Su et al. 1998; Hrivnak et al. 1999; Meixner et al. 1999; Sahai et al. 1999; Kwok et al. 2000; Ueta et al. 2000; Hrivnak et al. 2001; Gledhill 2005; Si\'odmiak et al. 2008; Lagadec et al. 2011a), with the goal of better understanding the shaping mechanisms of PPNe and PNe. The central stars of dusty pre-PNe are invariably obscured making their identification difficult, and in contrast the dust shell is subordinate to the central star in many high-latitude objects. This variety of properties has led to a myriad of different identification schemes and criteria in their identification to date. These include the presence of an infrared (IR) excess (Zuckerman 1978; Parthasarathy \& Pottasch 1986; Pottasch \& Parthasarathy 1988; Hrivnak, Kwok \& Volk 1989), including the utilization of IRAS colour-colour diagrams (e.g. Van der Veen \& Habing 1988; Van der Veen, Habing \& Geballe 1989; Preite-Martinez 1988; Manchado et al. 1989; Oudmaijer et al. 1992; Hu et al. 1993; Garc\'ia-Lario et al. 1997; Sahai et al. 2007). Identifying PAGB stars on the basis of their mid-IR spectra is also undertaken, including the presence of the distinctive $21~\micron$ feature in carbon-rich pre-PNe (Kwok, Volk \& Hrivnak 1989; Cerrigone et al. 2011). More recently, surveys for new PAGB stars using near-IR photometric data (Ramos-Larios et al. 2009, 2012), and the related R\,CrB stars using either near- and mid-IR colours (Tisserand et al. 2011; Tisserand 2012), or ASAS-3 optical light curves (Tisserand et al. 2013), have been undertaken. An improved understanding of this evolutionary phase is dependent on determining accurate distances to a large sample of objects, which can be used to furnish meaningful physical characteristics. Unfortunately at present, reliable distances have so far been determined for only a small fraction of these objects. It is therefore imperative that an accurate method for calculating distances to PAGB objects is determined. Since the PAGB phase is characterised by a near constant bolometric luminosity for their central stars, from the AGB-tip to the beginning of the white-dwarf cooling track (Paczy\'nski 1971; Sch\"onberner 1983; VW94; Bl\"ocker 1995), we can use a standard-candle luminosity to estimate the distances to them. This is the main focus and legacy of this paper, described in detail in Section\,\ref{Method}, below. %*********************************************************************** \subsection{Nomenclature}\label{Nomenclature} %*********************************************************************** For the benefit of the reader, we briefly describe the nomenclature that we have adopted in this paper. The generic term, PAGB star, includes all objects evolving from the AGB to the beginning of the white dwarf cooling track, with or without a surrounding nebula, though in practice, planetary nebulae (PNe) are defined separately (Kwok 1993, 2010; Frew \& Parker 2010). With a few exceptions (e.g. Jacoby et al. 1997; Alves, Bond \& Livio 2000), most Population\,II PAGB stars do not have extensive surrounding dust shells nor ionized PNe. We use the term pre-planetary nebulae (PPNe) to describe the non-ionized, dusty nebulae that scatter the light of their embedded central stars, and which emit in the thermal-IR (van der Veen \& Habing 1989; Pottasch \& Parthasarathy 1988). Spectra of their central stars range from B-type at the hot end to late-K or even early M-type at the cool end (Volk \& Kwok 1989; van Hoof, Oudmaijer \& Waters 1997; Van Winckel 2003). %check Frosty Leo Once the effective temperature of the central star reaches about 20\,kK, the surrounding material is photoionized to produce a young PN. The term ``transition object'' is sometimes used to describe objects just commencing this process of ionization (Su\'arez et al. 2006; Cerrigone et al. 2008; Frew, Boji{\v c}i{\'c} \& Parker 2013), graphically demonstrated by the recent evolution of CRL\,618 (Tafoya et al. 2013). Ueta, Meixner \& Bobrowsky (2000) classified resolved PPNe into two groups: the Star-Obvious Low-level Elongated (SOLE) nebulae, which have a visible nebula around an obvious central star, and the DUst-Prominent Longitudinally EXtended (DUPLEX) objects, which typically have faint or even invisible central stars obscured by a dusty torus. These differences extend to other observed properties. In general the SOLE objects have bluer infrared colours than the dustier DUPLEX nebulae (Ueta et al. 2000; Si\'odmiak et al. 2008). Similarly the SEDs are different: the SEDs of the SOLE objects are typically two peaked, dominated by the stellar photosphere and the dust component, while the DUPLEX sources have a very prominent dust peak, with little or no optical peak (Si\'odmiak et al. 2008). Si\'odmiak et al. (2008) also noted the different distributions in Galactic latitude of the two classes, and along with Meixner et al. (2002), suggested that the DUPLEX nebulae derive from more massive progenitor stars and are the natural precursors to bipolar PNe, which have been shown to have a lower scale-height than other PNe (Corradi \& Schwarz 1995; Phillips 2001; Frew 2008). We will investigate this problem in more detail in a later paper in this series. It has become apparent that many PAGB stars have little or no dust around them, and these are generally thought to be of low mass (e.g. Alcolea \& Bujarrabal 1991; Bujarrabal et al. 2013). Note that the ``\emph{likely}'' PAGB section of the Toru\'n catalogue includes the UU\,Her variables (e.g. Sasselov 1984), named after the prototype UU\,Herculis, whose own classification is debatable (Klochkova et al. 1997). We consider these stars to be the lower-luminosity halo analogues of the old-disk ``89~Herculis'' stars (Gillett, Hyland \& Stein 1970; Bujarrabal et al. 2007, and references therein). %%TO CHECK THESE REFS AGAIN Several other groups of uncommon stars are often classed as possible PAGB objects. These include the RV\,Tauri stars (Preston et al. 1963; Goldsmith et al. 1987; Van Winckel et al. 1999), pulsating yellow supergiants related to the Type II cepheids (Wallerstein 2002). The more luminous RV\,Tau stars are usually considered to be PAGB stars with low initial masses (Jura 1986; cf. Matsuura et al. 2002). For these stars, distances can be estimated using the period-luminosity (P-L) relation for Population II Cepheids (Alcock et al. 1998; Matsunaga et al. 2006; Soszy\'nski et al. 2008; Matsunaga, Feast \& Menzies 2009).\footnote{The less-luminous Type\,II Cepheids (W\,Vir and BL\,Her stars) are probably in an intermediate evolutionary phase between the blue horizontal branch and the base of the AGB, or are on a blue loop from the lower AGB (e.g. Maas, Giridhar \& Lambert 2007). They are not considered further.} A detailed study of their distances and space distribution will be the subject of the second paper in this series (Vickers et al., in preparation). The hydrogen-deficient R\,Coronae Borealis stars (Clayton 1996, 2012), and their hotter kin, the extreme helium stars (e.g. Pandey et al. 2001; Jeffery 2008) will also be evaluated in that work. %*********************************************************************** \subsection{Scientific Motivation }\label{Motivation} %*********************************************************************** The compilation of the Toru\'n Catalogue of Galactic PAGB and related objects (Szczerba et al. 2007, 2012) now provides a central repository of information for all currently identified Galactic PAGB stars, facilitating a wider study of these objects. To date the general physical characteristics of PAGB objects have only been determined from relatively small samples, using well studied PAGB objects with ample data available. Prior to the Toru\'n catalogue, it was necessary to collect scattered photometric and spectroscopic data to find candidate PAGB stars; i.e. a source displaying canonical PAGB colours (van der Veen \& Habing 1989; Pottasch \& Parthasarathy 1988). This effectively meant that that a large-scale investigation of the Galactic PAGB population was not feasible. While some of the the data sets available in the Toru\'n catalogue are limited in quality, more recent all-sky surveys such as the AKARI (Astro-F) survey (Ishihara et al. 2010) and Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) survey provide photometric data that is more sensitive and of higher resolution than those previously utilised. % in this technique (e.g. de Ruyter et al. 2005, 2006; Sahai et al. 2007). Here we present, for the first time, a homogenised catalogue of distances of all known Galactic PAGB objects available in the Toru\'n catalogue. Distances have been calculated using the observed SEDs, generated using the photometric and spectroscopic data gathered in the Toru\'n catalogue, as well as additional photometric data from recent all-sky surveys. Due to the narrow distribution of white dwarf masses (Vennes et al. 2002; Kleinman et al. 2013; and others), we have adopted assumed luminosities for specific sub-types rather than attempting to determine an individual luminosity for each object (see Section\,\ref{sec:luminosity}, below). Our distance catalogue will allow new insights into this brief, poorly understood phase of late-stage stellar evolution by allowing a population study based on improved distance estimates. The paper will proceed as follows: in \textsection~\ref{Material} we outline the material used including additional data sources taken from the literature, and in \textsection~\ref{Method} we detail the method used for the SED derived distances. In \textsection~\ref{Results} we provide the reader with a sample table of the SED calculated distances, and a comparison of the results with independent literature distances. In \textsection~\ref{Summary} we summarise our findings and give suggestions for future work. The resulting catalogue of distances as well as the fitted SEDs will be available in full as an online supplement. %*********************************************************************** | \label{Summary} %************************************************************************************ The compilation of the Toru\'n catalogue (Szczerba et al. 2007, 2012) gathered a wide assortment of flux data for all Galactic PAGB objects known at that time. We have used these data, adding more recent fluxes from the literature to build an homogenised set of distances by modelling their observed SEDs with one or more black body curves. Total fluxes were calculated for each object by numerically integrating the fitted curves. In a follow up paper we investigate the potential discrepancy in the integrated fluxes when using model atmospheres to model the central star as opposed to black body fitting. We expect that for severely reddened stars and stars with a high $F_{\rm IR}/F_\star$ model atmospheres will make little difference compared to the UV-bright stars where we expect a larger difference in integrated fluxes. The assumed luminosity (in solar units) was derived using the empirical core-mass luminosity relation for PAGB evolution from VW94, and a set of criteria separating the different populations of PAGB objects. Distances were computed by equating an assumed luminosity with the integrated flux of each source thus creating a homogenised set of distances, presented in full as an online supplement. The calculated distances were compared to several independently derived literature values measured using a variety of methods, in order to ascertain the accuracy of our approach. In Section~\ref{sec:Lit_comp}, we showed that our derived distances are in good agreement with a range of literature values. In this way we have effectively demonstrated that the SED technique is a valid method for calculating statistical distances to PAGB and related objects. In a follow-up paper (Vickers et al., in preparation), we will determine distances to the remaining objects in the Toru\'n Catalogue, namely the RV\,Tauri stars, using instead empirical period-luminosity relations, as well as the R\,CrB stars (and related hydrogen-deficient objects). In a further paper, we will investigate the population characteristics of a relatively complete volume-limited sample of Galactic disk PAGB objects for the first time. Such a census can be used for understanding the population demographics of Galactic PAGB objects, and their relationships with their precursor AGB stars and descendent PNe (Frew \& Parker 2006; Frew 2008; Frew et al. 2014b). Volume-limited samples are an under-appreciated tool for studying stellar populations (Frew \& Parker 2012), having the power to unlock the vital characteristics of Galactic PAGB objects, needed to understand the possible shaping mechanisms of their progenitors (Balick \& Frank 2002). Specifically we will endeavour to determine the scale heights (and hence progenitor ages and masses) of the various subgroups of PAGB stars and relate these to the morphological and dust properties of the resolved nebulae (Ueta et al. 2000; Si\'odmiak et al. 2008), and those objects that possess Keplerian dust disks (see Van Winckel et al. 2006; Hinkle et al. 2007; van Aarle et al. 2011; Acke et al. 2013). Our upcoming analysis will be undertaken with much larger samples than have been utilised previously (Likkel, te Lintel Hekkert \& Chapman 1993). Finally we expect the data avalanche from modern multi-wavelength surveys to aid in the discovery of many more PAGB stars and PPNe in the Galaxy. These will be incorporated into our new relational database of PNe and PAGB stars currently under construction at Macquarie University, in conjunction with the CDS, Strasbourg (Boji{\v c}i{\'c} et al., in preparation). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % In high-mass AGB stars, 13C can be produced by CNO-cycle hot-bottom burning of the third dredge-up 12C and then transported to the photosphere. 13C is readily synthesized in high mass AGB stars (Lattanzio \& Wood 2003). A low 12C/13C ratio is expected when a star undergoes hot-bottom burning. The lower mass limit for this process to occur is > 4 M? (Sackmann \& Boothroyd 1992). The implication of our derived 12C/13C ratio is that these extreme OH/IR stars have originated from the high mass end of intermediate-mass stars %OH 26.5+0.6, OH 127.8+0.0, OH 30.1-0.7, AFGL 5379, WX Psc %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %*********************************************************************** | 14 | 3 | 1403.7230 |
1403 | 1403.0860_arXiv.txt | Recently Strohmayer and Mahmoodifar presented evidence for a coherent oscillation in the X-ray light curve of the accreting millisecond pulsar XTE~J1751$-$305, using data taken by RXTE during the 2002 outburst of this source. They noted that a possible explanation includes the excitation of a non-radial oscillation mode of the neutron star, either in the form of a g-mode or an r-mode. The r-mode interpretation has connections with proposed spin-evolution scenarios for systems such as XTE~J1751$-$305. Here we examine in detail this interesting possible interpretation. Using the ratio of the observed oscillation frequency to the star's spin frequency, we derive an approximate neutron star mass-radius relation which yields reasonable values for the mass over the range of expected stellar radius (as constrained by observations of radius-expansion burst sources). However, we argue that the large mode amplitude suggested by the Strohmayer and Mahmoodifar analysis would inevitably lead to a large spin-down of the star, inconsistent with its observed spin evolution, regardless of whether the r-mode itself is in a stable or unstable regime. We therefore conclude that the r-mode interpretation of the observed oscillation is not consistent with our current understanding of neutron star dynamics and must be considered unlikely. Finally we note that, subject to the availability of a sufficiently accurate timing model, a direct gravitational-wave search may be able to confirm or reject an r-mode interpretation unambiguously, should such an event, with a similar inferred mode amplitude, recur during the Advanced detector era. | With improved observational sensitivity in both the electromagnetic and gravitational-wave channels, we are approaching the time when we may be able to probe neutron star interiors through asteroseismology. This era may, in fact, already have begun with observations of quasiperiodic oscillations in the tails of magnetar giant flares (see, e.g. \citealt{israel05}, \citealt{SW06}) and the suggestion that the observed oscillations can be mapped to shear modes of the star's crust (e.g. \citealt{SA07}). However, the main lesson learned from this exercise is that the asteroseismology programme is complicated, owing in part to the unknown nature of the state of matter inside the star and the form of the internal magnetic field \citep{toypaper, gabler13}. In order to progress further, we need theoretical models that account for as much of the involved complex physics as possible (or, indeed, palatable because the problem becomes increasingly messy). We also need more observational data. In this context, it is interesting to note the recent suggestion by \shortcite{strohmayermahmoodifar13} that a non-radial oscillation mode may have been present in the X-ray data associated with the 2002 discovery outburst of XTE~J1751$-$305, an accreting neutron star in a low-mass X-ray binary (LMXB) with spin frequency $\nu= 435$~Hz. The observed frequency $0.5727597\times \nu$ is tantalizingly close to the frequency of the star's quadrupole r-mode. In Newtonian gravity, this r-mode would have frequency $2\nu/3$ to leading order of a slow-rotation expansion. Noting this proximity, \shortcite{strohmayermahmoodifar13} discuss whether the observations could be consistent with an r-mode excited in the bulk of the star. This would be an exciting result, given the suggestion that gravitational-wave emission may drive this particular mode unstable at some critical rotation rate, providing a mechanism to prevent further spin-up by accretion (\citealt{anderssonetal99}; see also \citealt{Ho,haskelletal12}, for a recent assessment of this idea). At first sight, the r-mode interpretation would seem unlikely, because the observed frequency is close to the mode frequency in a frame co-rotating with the star whereas one might expect a distant observer to see the inertial frame mode frequency, which is $4\nu/3$ to leading order. The first part of any r-mode interpretation has to deal with this issue. The required explanation was provided by \shortcite{NL}, who demonstrated that a nonradial oscillation mode can indeed lead to modulations of an X-ray hotspot being observed at the rotating frame mode frequency. This conclusion is key to the r-mode interpretation of the XTE~J1751$-$305 data. However, this is not sufficient to make the connection credible. Strohmayer and Mahmoodifar concluded in part that to interpret the observed oscillation frequency as an r-mode, it was necessary for the star's spin rate to be such that it was close to a resonance between a crustal torsional mode and the r-mode frequency. Furthermore, for this to be the case, they noted that the star's crust must have a shear modulus approximately twice as large as is generally predicted (cf. the results of \citealt{shearmod}). One might not be comfortable with such a sizeable change of a relatively well understood part of the star's physics, but it is a possible explanation. However, it turns out that one does not need to tweak crust physics to explain the observed frequency in terms of a global r-mode. A large correction to the r-mode frequency is caused by the combined relativistic effects of the gravitational redshift and rotational frame-dragging. These effects were calculated some time ago by \citet{LFA03}. According to their analytic expression for a uniform density barotropic model (equation~36 of \citealt{LFA03}) the quadrupole ($m=2$) r-mode should have a co-rotating r-mode frequency $\kappa\times\nu$ where \begin{equation} \kappa = {2 \over 3} \left[ 1 - {8\over 15} \left( {M\over R}\right) \right] , \label{pNfreq}\end{equation} to first post-Newtonian (1PN) order. This shows that, for a typical neutron star compactness of $M/R\approx 0.2$, the relativistic corrections is at the 10\% level. This is significantly larger than the rotational and crust corrections (unless one invokes resonances) considered in \citet{strohmayermahmoodifar13}. Moreover, the relativistic effects tend to \emph{lower} the (rotating frame) mode frequency, exactly as required to explain the observations, which require $\kappa \approx 0.57$. This motivates us to take a closer look at the problem. In Section~\ref{mass} we go beyond the rough estimate of equation (\eqref{pNfreq}) and argue that we can use the observational results to constrain the relationship between mass and radius for this neutron star. If one additionally assumes that the radius lies in the range inferred from radius-expansion burst sources \citep{LS}, estimates for the mass itself are obtained. We show that this procedure yields a sensible result. However, this is only part of the story---it is also necessary to explain the observed variation in spin frequency for this system. In Section \ref{observations} we first consider the issue of spin evolution including only standard accretion torques and magnetic fields, without any r-mode excitation. We show that the standard models fail to explain the observed frequency variation of XTE~J1751$-$305, as different estimates/constraints on the star's magnetic field fail to agree; we find the same disagreement for another LMXB, IGR~J00291+5934. In Section \ref{scenarios} we consider the effect of adding an r-mode into the model, with the large amplitude indicated by the Strohmayer \& Mahmoodifar analysis. We show that, in all plausible scenarios (i.e., stable and unstable and unsaturated and saturated), the inclusion of the r-mode over the duration of the X-ray outburst (10 days or so) in fact makes the observed spin variation much \emph{more} difficult to explain (Strohmayer \& Mahmoodifar noted this problem while considering only the unstable scenario). We therefore conclude that the r-mode interpretation of the data is hard to sustain, but note that future gravitational wave detectors could, if a sufficiently accurate timing model is available, be able to directly test if an r-mode is present at a similar level in any future outburst of XTE~J1751$-$305 or a similar system. | \label{discussion} We showed that the observed frequency of the oscillation seen in XTE~J1751$-$305 during the 2002 outburst, if interpreted as an r-mode, would correspond to sensible values of mass and radius for the neutron star. However, we also argued that the presence of an r-mode with the large amplitude that the observations suggest, cannot be reconciled with the observed spin-up observed over the duration of the outburst. We conclude that the presence of an r-mode in XTE~J1751$-$305 during the 2002 outburst is not consistent with our current understanding of neutron star dynamics and must therefore be considered unlikely. Nevertheless, one can look to gravitational-wave observations to see if they might be able to rule out (or indeed support) the r-mode scenario for similar events in the future. For the purpose of such a search in the case of XTE~J1751$-$305, it is interesting to note that the frequency of the signal would be known. The observed electromagnetic signal would carry the rotating frame frequency, but a gravitational-wave counterpart would be found at the inertial frame frequency. This means that one ought to search at \begin{equation} \nu_\mathrm{GW} \approx \left| \kappa - m \right| \nu \approx 621\ \mathrm{Hz}. \end{equation} Moreover, it is easy to make a back-of-the-envelope estimate of the effective gravitational amplitude. The associated strain can be obtained from \begin{equation} h_0 = \sqrt{{8\pi\over 5}} {G \over c^5} {\alpha \over d} \omega_{\rm GW}^3 MR^3 \tilde J, \end{equation} where we take a conservative distance to the source of $d=9$~kpc and $\omega_{\rm GW} = 2\pi \nu_{\rm GW}$ \citep{owen2010}. Inserting a mode amplitude of $\alpha \sim 10^{-3}$, this leads to (for the preferred values of mass and radius from Section~\ref{mass}) \begin{equation} h_0 \approx 1 \times 10^{-24} \alpha_{-3}. \end{equation} Assuming that the mode remains at this amplitude over the outburst, this would be a periodic signal, the detectability of which increases with the observation time. The effective amplitude is then \begin{equation} h_\mathrm{eff} = h_0 \sqrt{ T} \approx 9 \times 10^{-22} \alpha_{-3} \left(\frac{\Delta T_{\rm obs}}{6 \, \rm days}\right)^{1/2} {\,\, \rm Hz}^{-1/2} . \end{equation} The outburst in XTE~J1751$-$305 in fact preceded the first LIGO science run (S1) by a matter of a few months, so there is no possibility of carrying out a gravitational-wave search on archival LIGO science data. However, we note that this would have corresponded to a signal-to-noise of about $10$ if present in LIGO S5 data, and about $100$ in Advanced LIGO data, \emph{assuming that one can perform a fully phase coherent search over the observation period}. \citet{strohmayermahmoodifar13} did indeed find that the observed oscillation was highly coherent, quoting the oscillation frequency (in Hertz) to the sixth decimal place, i.e. at the resolution of the Fourier transform $\sim 1 /$($6$ days). See also their Figure 4. Clearly, the oscillation itself was intrinsically stable, and also the orbital solution was known sufficiently accurately to allow the effects of Doppler modulation due to the orbit to be largely removed over the observation interval. Should such an accurate timing solution be available for any future outburst, a fully coherent gravitational wave search could be carried out. Alternatively, in the absence of such an accurate timing solution, a search over a parameter space that allows for the orbital uncertainties could be carried out, but this would increase the signal-to-noise required to claim detection, as discussed in detail in \citet{wetal08}. Clearly, the signal analysis procedures for carrying out such a search merit further investigation and could be put to use if this (or a similar accreting system) were to display similar behaviour during a future science run of an advanced gravitational-wave detector. | 14 | 3 | 1403.0860 |
1403 | 1403.5168_arXiv.txt | \label{sec:intro} Large-scale magnetic fields exist in cosmic structures such as galaxies and galaxy clusters, however the origin of the fields remain unexplained. Recent gamma ray observations suggest the existence of magnetic fields even in void regions~\cite{Tavecchio:2010mk,Neronov:1900zz,Ando:2010rb,Taylor:2011bn,Takahashi:2013uoa,Finke:2013bua}, and the amplitude of such intergalactic magnetic fields were derived to be stronger than $\sim 10^{-15}$G. Although this lower bound has astrophysical uncertainties (see e.g.~\cite{Broderick:2011av,Miniati:2012ge}), such large-scale magnetic fields hint that the magnetic fields are of cosmological origin. Here it should be noted that the equations of motion of the electromagnetic fields in a Friedmann universe reduce to those in Minkowski space, since the standard Maxwell theory is conformally invariant. Therefore the conformal invariance should be broken for the electromagnetic fields to be significantly produced by the cosmological background. Such mechanisms of cosmological magnetogenesis have been embedded in the inflationary universe in various works, e.g.,~\cite{Turner:1987bw,Ratra:1991bn,Garretson:1992vt,Gasperini:1995dh,Giovannini:2000dj,Davis:2000zp,Bamba:2003av,Martin:2007ue,Demozzi:2009fu,Kanno:2009ei}.\footnote{As an alternative way of cosmologically producing magnetic fields, \cite{Vachaspati:1991nm,Cornwall:1997ms,Vachaspati:2008pi} have studied magnetogenesis during phase transitions.} However the studies have revealed that inflationary magnetogenesis cannot produce significant magnetic fields without running into at least one of the inconsistencies: (i) the electromagnetic fields obtain strong couplings and therefore the theory becomes uncontrollable~\cite{Gasperini:1995dh,Demozzi:2009fu,Fujita:2012rb}, (ii) too large electric fields are produced such that their backreaction can spoil inflation and/or magnetogenesis~\cite{Bamba:2003av,Demozzi:2009fu,Kanno:2009ei,Fujita:2012rb}. In particular, the backreaction of the electric fields is a problem inherent to inflationary magnetogenesis; after magnetogenesis the magnetic fields decay as $B \propto a^{-2}$ where $a$~is the scale factor, therefore the present day magnetic fields should be remnants of substantial photon production during a short period of time in the very early universe. In other words, the electromagnetic vector potential during inflation is required to possess a large time-derivative, and therefore leads to excessive production of electric fields. Recently it has also been pointed out that the generated electromagnetic fields can induce cosmological density perturbations beyond the observed value, and works~\cite{Suyama:2012wh,Giovannini:2013rme,Ringeval:2013hfa,Nurmi:2013gpa,Fujita:2014sna} have imposed further constraints on inflationary magnetogenesis. (See also~\cite{Shaw:2010ea,Yamazaki:2012pg,Camera:2013fva,Berger:2014wta} and references therein for constraints on primordial magnetic fields from CMB and large-scale structure data.) However, we stress that the inflationary expansion itself is not a necessary condition for magnetogenesis; it is the breaking of the conformal invariance that allows magnetic field production by the cosmological background. In this paper, we investigate post-inflationary magnetogenesis by breaking the conformal invariance of the Maxwell theory after inflation. The conductivity of the universe becomes high during reheating and thereafter the magnetic fields decay as $B \propto a^{-2}$~\cite{Turner:1987bw}. Therefore we focus on the epoch between the end of inflation and (p)reheating, during which the inflaton field oscillates around its potential minimum and the universe is effectively dominated by cold matter. The conformal invariance of the Maxwell theory is broken by couplings between the electromagnetic fields and scalar degrees of freedom which can be the inflaton or some other spectator field(s). By considering magnetogenesis both during and after inflation, we demonstrate that magnetic fields of~$10^{-15}$G or stronger can be produced at cosmological scales (say, on Mpc scales), without running into the strong coupling regime or producing too much electric fields. The problem of affecting the cosmological density perturbations is also ameliorated since the magnetic fields are enhanced after inflation. The proposed model is compatible even with high scale inflation. This paper is organized as follows. We first review the electromagnetic theory in an expanding universe in Section~\ref{sec:EMFRW}. Then we move on to study magnetogenesis during the inflationary epoch in Section~\ref{sec:during}, and post-inflationary magnetogenesis in Section~\ref{sec:after}. Here we will see that the magnetogenesis in each epoch alone are highly constrained by the strong coupling and backreaction problems. In particular, the constraints on inflationary and post-inflationary magnetogenesis will be translated into severe upper bounds on the inflation and reheating scales, respectively. Then in Section~\ref{sec:model} we discuss the combined scenario of inflationary {\it and} post-inflationary magnetogenesis. There we will see that such a two step model can overcome the challenges and efficiently produce large-scale magnetic fields. In Section~\ref{sec:couplings}, we present some examples of scalar couplings that break the conformal invariance of the Maxwell theory. Finally, we conclude in Section~\ref{sec:conc}. | \label{sec:conc} In this work, we explored cosmological magnetogenesis during the two phases when the universe is cold: the inflationary epoch, and the post-inflationary epoch prior to reheating, during which the universe is dominated by the oscillating inflaton. Magnetogenesis in each phase alone are highly constrained by the strong coupling and backreaction problems, however, we have found that the combined inflationary {\it and} post-inflationary magnetogenesis can overcome the difficulties and efficiently produce large-scale magnetic fields. In particular, we demonstrated that the combined inflationary/post-inflationary magnetogenesis scenario can produce magnetic fields stronger than $10^{-15}\, \mathrm{G}$ on Mpc scales without running into the strong coupling regime, or producing too large electric fields that would dominate the universe. The proposed model is compatible even with high scale inflation. The strong enhancement of the magnetic fields is made possible in the two step scenario due to the magnetogenesis in the two epochs working in very different ways. The magnetic enhancement in the post-inflationary universe reduces the need for a significant production of magnetic fields during inflation, and thus relaxes the constraints on inflationary magnetogenesis, including those from the backreaction problem and the excessive production of cosmological density perturbations. On the other hand, inflationary magnetogenesis enhances large-scale magnetic fields and thus can relatively suppress the small-scale (but super-horizon) electromagnetic fields produced during post-inflationary magnetogenesis. Therefore the two phases of magnetogenesis mutually prevent the electromagnetic fields from dominating the universe in each epoch, while maintaining magnetic power at large scales. Moreover, we have shown that the net magnetic enhancement from the two step magnetogenesis is much stronger than a naive product of the enhancements from the individual phases of magnetogenesis. In order to generate large magnetic fields from the cosmological background, we considered breaking the conformal invariance of the Maxwell theory $-I^2 F_{\mu \nu} F^{\mu \nu }/4$ through a time dependent coupling~$I$ that scales as a power-law of the scale factor. As we discussed in Section~\ref{sec:couplings}, such a dynamical~$I$ can arise from couplings with the inflaton or some other spectator field(s) that rolls along its effective potential, or oscillates about the potential minimum. Here we note that the produced electromagnetic fields may backreact on the directly coupled scalars in some cases. To see whether such effects become important requires a detailed analysis of explicitly constructed models. We leave this for future work. The difficulty of post-inflationary magnetogenesis alone was presented in the form of the strict upper bound on the reheating scale $H_{\mathrm{reh}} \ll 10^{-23} \, \mathrm{MeV}$ (\ref{Hrehbound}), which is incompatible with BBN. Upon deriving this bound we have assumed the scaling~$I \propto a^{-n}$, and so the generality of the reheating bound remains an open question. For example, electromagnetic couplings with oscillating scalars can lead to further amplification of the magnetic fields through parametric resonance, which may allow the post-inflationary universe alone to create large magnetic fields, without the aid of inflationary magnetogenesis. Thus it would be interesting to derive a generic bound on post-inflationary magnetogenesis without specifying how the electromagnetic mode function evolves in time. Such an analysis was carried out for inflationary magnetogenesis in~\cite{Fujita:2012rb}, and it may be possible to apply their discussions to the post-inflationary universe as well. We also note that the magnetic enhancement in the post-inflation universe should be affected by the details of the reheating process. We have assumed the conductivity of the universe to suddenly become high towards the end of the inflaton-dominated epoch, however the conductivity may start to gradually increase from earlier times if the oscillating inflaton decays perturbatively. On the other hand, preheating~\cite{Kofman:1994rk,Kofman:1997yn} may lead to a sudden growth of the conductivity. Although more detailed work will be required to verify whether there actually exists intergalactic magnetic fields, investigation of cosmological magnetic fields may provide an observational window into the very early universe. We demonstrated that large-scale magnetic fields can actually be created from the cosmological background. We hope that our mechanism will provide new insights into explaining our magnetized universe from the cosmological point of view. | 14 | 3 | 1403.5168 |
|
1403 | 1403.5732_arXiv.txt | We present new constraints on the spectral index $n_T$ of tensor fluctuations from the recent data obtained by the BICEP2 experiment. We found that the BICEP2 data alone slightly prefers a positive, "blue", spectral index with $n_T=1.36\pm0.83$ at $68 \%$ c.l.. However, when a TT prior on the tensor amplitude coming from temperature anisotropy measurements is assumed we get $n_T=1.67\pm0.53$ at $68 \%$ c.l., ruling out a scale invariant $n_T=0$ spectrum at more than three standard deviations. These results are at odds with current bounds on the tensor spectral index coming from pulsar timing, Big Bang Nucleosynthesis, and direct measurements from the LIGO experiment. Considering only the possibility of a "red", $n_T<0$ spectral index we obtain the lower limit $n_T > -0.76$ at $68 \%$ c.l. ($n_T>-0.09$ when a TT prior is included). | The recent detection of B-mode polarization made by the BICEP2 experiment \cite{bicep2} clearly represents one of the major discovery in cosmology in the past twenty years. While the BICEP2 result clearly needs to be confirmed by future experiments, it is timely and important to fully analyze the BICEP2 data and to identify all possible inconsistencies at the theoretical level. In this brief note we focus our attention on the spectral index of tensor fluctuations $n_T$. Indeed, a crucial prediction of inflation is the production of a stochastic background of gravity waves (\cite{GWs}) with a slightly tilted spectrum, \begin{equation} n_T = -2\epsilon ~, \end{equation} \noindent where $\epsilon=-\dot{H}/H^2$ denotes a slow roll parameter from inflation ($H$ is the Hubble rate during the inflationary stage). In standard inflation $\epsilon$ is strictly positive~\cite{lr} and in the usual parameter estimation routines, the tensor spectral index is assumed to be ``red'', or negligible. However, in recent years, a set of inflationary models has been elaborated where the spectral index of tensor modes could be positive, $n_T>0$, i.e. ``blue''. A first attempt to compare these models with observational data has been made in \cite{camerini}. The main theoretical problem for the production of a blue spectrum of gravitational waves (BGW) is that the stress-energy tensor must violate the so-called Null Energy Condition (NEC). In a spatially flat FRW metric, a violation of NEC indeed corresponds to the inequality $\dot{H}<0$ and is ultimately the reason for the red tensor spectrum in standard inflation. Models that violates NEC have been already presented. For example, in the so-called super-inflation models \cite{super} where inflation is driven by a component violating the NEC a BGW spectrum is expected. Models based on string gas cosmology as in \cite{stringgas}, where scalar metric perturbations are thought to originate from initial string thermodynamic fluctuations \cite{stringpert}, also can explain a BGW background. A BGW spectrum is also a generic prediction of a class of four-dimensional models with a bouncing phase of the universe \cite{bounce}. To induce the bounce, the stress-energy tensor must violate the null energy condition (NEC). G-inflation \cite{yoko}, has a Galileon-like nonlinear derivative interaction in the Lagrangian with the resultant equations of motion being of second order. In this model, violation of the null energy condition can occur and the spectral index of tensor modes can be blue. BGW may also be present in scalar-tensor theories and $f(R)$ gravity theories. It is therefore timely to investigate the constraints on the tensor spectral index $n_T$ from the BICEP2 data. Strangely enough, no constraint on this parameter has been presented by the BICEP2 collaboration while, as we discuss in the next section, we found that the BICEP2 data could provide interesting results on this parameter. | In this brief note we have presented new constraints on the spectral index $n_T$ of tensor fluctuations from the recent data obtained by the BICEP2 experiment. We found that the BICEP2 data alone slightly prefers a positive, "blue", spectral index with $n_T=1.36\pm0.83$ at $68 \%$ c.l.. However, when a TT prior on the tensor amplitude coming from temperature anisotropy measurements is assumed we get $n_T=1.67\pm0.53$ at $68 \%$ c.l., ruling out a scale invariant $n_T=0$ spectrum at more than three standard deviations. Considering only the possibility of a "red", $n_T<0$ spectral index we obtain the lower limit $n_T > -0.76$ at $68 \%$ c.l. ($n_T>-0.09$ when a TT prior is included). These results are at odds with current upper limits on the tensor spectral index coming from observations of pulsar timing, Big Bang Nucleosynthesis, and from direct upper limits from the LIGO experiment (see e.g. \cite{upperlimitsnt}). Considering $r_{0.002}=0.2$ and using the method adopted in \cite{upperlimitsnt} we found the current upper limits on $n_T$: $n_T \le 0.81$, $n_T \le 0.29$ and $n_T \le 0.15$ at $68 \%$ c.l. from pulsar timing, LIGO and BBN respectively. The LIGO and BBN limits are in strong tension with the BICEP2+CMB value. Therefore a positive spectral index does not provide an acceptable solution to the tension between the BICEP2 data and current upper limits on $r$ from temperature anisotropies. This indicates either the need of including extra parameters (as the running of the scalar spectral index \cite{bicep2} or extra neutrino species \cite{giusarma2014}) to relax current bounds on $r_{0.002}$ from temperature anisotropies or the presence of unresolved systematics in current CMB data. During the submission of this paper other works appeared discussing the possibility of a BGW from BICEP2 (see \cite{bgw}) but without presenting numerical constraints on $n_T$ and an independent analysis of the BICEP2 data. We also like to point out the discussion on the \texttt{cosmocoffee.info} website where results similar to ours have been presented by Antony Lewis. | 14 | 3 | 1403.5732 |
1403 | 1403.2438_arXiv.txt | A number of brown dwarfs are now known to be variable with observed amplitudes as large as 10--30\% at some wavelengths. While spatial inhomogeneities in cloud coverage and thickness are likely responsible for much of the observed variability, it is possible that some of the variations arise from atmospheric temperature fluctuations instead of, or in addition to, clouds. To better understand the role that thermal variability might play we present a case study of brown dwarf variability using a newly-developed one-dimensional, time-stepping model of atmospheric thermal structure. We focus on the effects of thermal perturbations, intentionally simplifying the problem through omission of clouds and atmospheric circulation. Model results demonstrate that thermal perturbations occurring deep in the atmosphere (at pressures greater than 10~bar) of a model T-dwarf can be communicated to the upper atmosphere through radiative heating via the windows in near-infrared water opacity. The response time depends on where in the atmosphere a thermal perturbation is introduced. We show that, for certain periodic perturbations, the emission spectrum can have complex, time- and wavelength-dependent behaviors, including phase shifts in times of maximum flux observed at different wavelengths. Since different wavelengths probe different levels in the atmosphere, these variations track a wavelength-dependent set of radiative exchanges happening between different atmospheric levels as a perturbation evolves in time. We conclude that thermal---as well as cloud---fluctuations must be considered as possible contributors to the observed brown dwarf variability. | After more than a decade and a half of surveys for brown dwarf variability we now know that the emergent spectrum of many L- and T- type brown dwarfs indeed varies with time \citep[e.g.,][]{tinney&trolley1999,bailerjones&mundt1999,artigauetal2009,radiganetal2012,buenzlietal2013}. Broadband, near-infrared flux variations can be as large as 10--30\%, can occur on timescales from 1--100~hours, and can be non-periodic \citep{bailerjones&mundt2001a,gelinoetal2002,artigauetal2009,radiganetal2012,gillonetal2013}. Spectroscopic and multi-band photometric studies have revealed complex, wavelength-dependent lightcurves \citep{radiganetal2012,buenzlietal2013}. In some cases, spectra show periodic brightness fluctuations with wavelength-dependent phase lags, which can be as large as 180$^{\circ}$ (i.e., shifted by half of a cycle) \citep{buenzlietal2012}. Clouds sculpt the emergent spectra of essentially all spectral classes of brown dwarfs, although their impact is most notable in the L-dwarfs \citep{leggettetal1998,chabrieretal2000,allardetal2001,ackerman&marley2001,tsuji2002, golimowskietal2004,knappetal2004,burrowsetal2006,stephensetal2009,morleyetal2012}. Given the strong evidence for the presence of clouds in brown dwarf atmospheres, and the ability of a continuum opacity source to limit the depth of the wavelength-dependent photosphere \citep{ackerman&marley2001}, it is expected that these structures play some role in brightness variability \citep{radiganetal2012,apaietal2013}. Indeed thermal emission from the deep atmospheres of both Jupiter and Saturn is strongly modulated by cloud structures and Jupiter itself would show substantial variability if the disk were observed at 5~$\mu$m in integrated light \citep{gelino&marley2000}. However the intensity of radiation emitted by a planetary or brown dwarf atmosphere depends on many factors in addition to cloud structure. Atmospheric temperature and composition also control the thermal emission and it seems prudent to also consider the role such factors might contribute to variability. \citet{freytagetal2010} used a 2-D radiation hydrodynamics model to study atmospheric circulation and dust transport in M-dwarf and brown dwarf atmospheres. This work highlighted the importance of gravity waves and dust convection to maintaining clouds in brown dwarf atmospheres. They found that gravity waves are expected to be ubiquitous above the radiative-convective boundary and likely play an important role in cloud development and evolution. More recently, global 3-D, cloud-free models were used by \citet{showman&kaspi2013} to study large-scale flows and convection in the interiors and deep atmospheres of brown dwarfs. These models revealed that convection is strongly influenced by the relatively fast rotation rates of brown dwarfs and that thermal variations of order several Kelvin may be expected at the top of the model convective zone. This work also discussed a stratospheric circulation, driven by the interaction between atmospheric waves generated at the top of the convective zone and the mean stratospheric flow, that could lead to large ($\sim$50~K) temperature variations in the upper atmospheres of brown dwarfs. However, emission spectra were not computed, and so the influence of these variations on the brown dwarfs' spectra remain unclear. Here we provide a case study of the impact of atmospheric temperature fluctuations on the emission spectrum of a brown dwarf. We intentionally simplify the problem by neglecting clouds as well as chemical evolution, thus allowing us to explore the behaviors, timescales, and related spectral variability due to temperature fluctuations alone. Using a new 1-D, time-stepping radiative convective model for brown dwarfs, we first investigate the heating of an atmosphere due to an extended thermal pulse from deep within the convective zone. We then explore variability in the emission spectrum due to time-varying thermal fluctuations, introducing these perturbations at different atmospheric levels, and highlighting circumstances where the model can reproduce wavelength-dependent phase lags. | Using a one-dimensional model of brown dwarf atmospheric structure, we have studied the time-dependent evolution of the atmosphere in response to a variety of thermal perturbations. We omitted cloud and dynamical effects, choosing to concentrate on behaviors that arise strictly due to atmospheric thermal variations. Thermal perturbations of the deep atmosphere can be communicated to the upper atmosphere at shorter near-infrared wavelengths, although communication in the opposite direction is impeded by the lack of flux generated at these wavelengths by the relatively cool upper atmosphere. The response timescale of the atmosphere to thermal perturbations is typically 10--100~hours. Deep thermal perturbations can lead to brightness fluctuations at nearly all near-infrared wavelengths, and our model predicts that these could be observed on timescales of hundreds of hours. While it is not our goal to solve the entire problem of brown dwarf variability, our model can produce a number of the observed features, depending on the nature of the thermal perturbation. In the future, a full explanation of variability in brown dwarf thermal emission spectra must incorporate three-dimensional atmospheric and cloud dynamics, as well as the time-dependent evolution of thermal perturbations throughout the radiative portion of the atmosphere. | 14 | 3 | 1403.2438 |
1403 | 1403.0588.txt | {}{}{}{}{} % 5 {} token are mandatory \abstract % context heading (optional) % {} leave it empty if necessary {Virialized halos grow by the accretion of smaller ones in the cold dark matter scenario. The rate of accretion depends on the different properties of the host halo. Those halos for which this accretion rate was very fast and efficient resulted in systems dominated by a central galaxy surrounded by smaller galaxies at least two magnitude fainter. These galaxy systems are called fossil systems and they can be the fossil relics of ancient galaxy structures.} % aims heading (mandatory) {We started an extensive observational program to characterize a sample of 34 fossil group candidates spanning a broad range of physical properties.} % methods heading (mandatory) {Deep $r-$band images were obtained with the 2.5-m Isaac Newton Telescope and Nordic Optic Telescope. Optical spectroscopic observations were performed at the 3.5-m Telescopio Nazionale Galileo telescope for $\sim$ 1200 galaxies. This new dataset was completed with Sloan Digital Sky Survey Data Release 7 archival data to obtain robust cluster membership and global properties of each fossil group candidate. For each system, we recomputed the magnitude gaps between the two brightest galaxies ($\Delta m_{12}$) and the first and fourth ranked galaxies ($\Delta m_{14}$) within 0.5 $R_{{\rm 200}}$. We consider fossil systems those with $\Delta m_{12} \ge 2$ mag or $\Delta m_{14} \ge 2.5$ mag within the errors.} % results heading (mandatory) {We find that 15 candidates turned out to be fossil systems. Their observational properties are in agreement with those of non-fossil systems. Both follow the same correlations, but the fossil systems are always extreme cases. In particular, they host the brightest central galaxies and the fraction of total galaxy light enclosed in the brightest group galaxy is larger in fossil than in non-fossil systems. Finally, we confirm the existence of genuine fossil clusters.} {Combining our results with others in the literature, we favor the merging scenario in which fossil systems formed due to mergers of $L^\ast$ galaxies. The large magnitude gap is a consequence of the extreme merger ratio within fossil systems and therefore it is an evolutionary effect. Moreover, we suggest that at least one fossil group candidates in our sample could represent a transitional fossil stage. This system could have been fossil in the past, but not now due to the recent accretion of another group of galaxies.} | \label{intro} Fossil systems are group- or cluster-sized \citep{mendes06,cypriano06} objects whose luminosity is dominated by a very massive central galaxy. In the current cold dark matter (CDM) scenario, these objects formed hierarchically at an early epoch of the Universe and then slowly evolved until present day. That is the reason why they are called {\it fossils}. The study of this particular kind of objects started two decades ago, when \citet{ponman94} suggested that RX-J1340.6+4018 was probably the remains of an ancient group of galaxies. Later, \citet{jones03} gave the first observational definition of Fossil Groups (FGs) as systems characterized by a magnitude gap larger than 2 mag in the $r-$band between the two brightest galaxies of the system within half the virial radius. Moreover, the central galaxy should be surrounded by a diffuse X-ray halo, with a luminosity of at least $L_{\rm X} > 10^{42}$ $h_{50}^{-2}$ erg ${\rm s^{-1}}$, with the aim of excluding bright isolated galaxies. Many optical and X-ray observational properties of FGs have been studied, but always on small samples or individual systems. These properties can be grouped in: (i) properties of the intracluster hot component; (ii) properties of the galaxy population; and (iii) properties of the brightest group galaxy (hereafter BGG). Referring to the hot gas component, fossil and non-fossil systems generally show a similar $L_{\rm X}-T_{\rm X}$ relation \citep[see][]{khos07,harrison12}. Differences in scaling relations that combine both optical and X-ray properties were detected. In particular, some authors found different relations in optical vs X-ray luminosity ($L_{{\rm opt}}-L_{\rm X}$), X-ray luminosity vs velocity dispersion of the clusters galaxies ($L_{\rm X}-\sigma _{{\rm v}}$), and X-ray temperature vs velocity dispersion ($T_{\rm X}-\sigma _{{\rm v}}$). In these works, for any given L$_{{\rm opt}}$, FGs are more luminous and hotter in the X-rays than normal groups or clusters. These differences were interpreted as a deficit formation of $L^\ast$ galaxies in FGs \citep[see][]{proctor11}. In contrast, other authors such as \citet{voev10} and \citet{harrison12} did not find any different relation between X-ray and optical quantities for FGs and normal groups and clusters. They claimed that the previous differences were due to observational biases in the selection of FGs or inhomogeneity between the FGs and the comparison sample. In addition, high $S/N$ and high resolution X-ray observations of fossil systems seem to confirm that fossil systems are formed inside high centrally concentrated dark matter (DM) halos \citep{sun04,khos06}, with large mass-to-light ratios, which could indicate an early formation. Nevertheless, most of the fossil systems do not show cooling cores \citep[but see also][]{democles10} as normal clusters, suggesting that strong heating mechanisms, such as AGN feedback or cluster mergers, could heat the central regions of their DM halos \citep{sun04,khos04,khos06,mendes09}. The galaxy luminosity function (hereafter LF) is a powerful tool for studying the galaxy population in clusters. In the past, several works investigated the galaxy LF in fossil systems. They found that the LF of these objects are well fitted by single Schechter function, but there is a large variety of values in the faint-end slope ($\alpha$) of the LFs of FGs. In particular, the values of $\alpha$ measured goes from $-1.6$ to $-0.6$ \citep[see][]{cypriano06,khos06,mendes06,mendes09,zibetti09,aguerri11,lieder13}. Unfortunately, all these studies were performed on single FGs or very small samples, and a systematic study of LFs of statistically meaningful samples of FGs remains to be done. The brightest central galaxies of fossil systems are amongst the most massive and luminous galaxies known in the Universe. In fact, the luminosity and the fraction of light contained in the BGGs correlate with the magnitude gap \citep{harrison12}. Some observations \citep{khos06} show that these objects are different from both isolated elliptical galaxies and central galaxies in non-fossil clusters in the sense that they have disky isophotes in the centre and their luminosity correlates with velocity dispersion, while other authors \citep{labarbera09, jairo12} found no differences in isophotal shapes between fossil and non fossil central galaxies. In \citet{jairo12} we analised deep K-band images of 20 BGGs in fossil and non-fossil systems and showed that these galaxies follow the tilted fundamental plane of normal ellipticals \citep[see ][]{bernardi11}. This fact suggests that BGGs are dynamically relaxed systems that suffer dissipational mergers during their formation. On the other hand, they depart from both Faber-Jackson and luminosity-size relations. In particular, BGGs have larger effective radii and smaller velocity dispersions than those predicted by these relations. We infer that BGGs grew throughout dissipational mergers in an early stage of their evolution and then assembled the bulk of their mass through subsequent dry mergers. Nevertheless, stellar population studies of BGGs in fossil systems suggest that their age, metallicity and $\alpha$-enhancement are similar to those of bright ellipticals field galaxies \citep[see][]{labarbera09,paul13}. In numerical simulations, FGs are found to be a particular case of structure formation. They are supposed to be located in highly concentrated DM halos, so that they can assembly half of their DM mass at {\it z} $>$ 1. Then, the FGs grow via minor mergers only, and only accrete $\approx$ 1 galaxy from {\it z} $\approx 1$ down to present time, while regular groups accrete about three galaxies in the same range of time \citep{vonbenda08}. \citet{dariush07} show that the mass assembled at any redshift is higher in fossil than in non-fossil systems. This means that the formation time is, on average, shorter for FGs than for regular systems \citep{donghia05, vonbenda08}, leaving to FGs enough time to merge $L^\ast$ galaxies in one very massive central object. In fact, simulations predict that the timescale for merging via dynamical friction is inversely proportional to the mass of the galaxy, thus favouring the merging of larger objects. So, the dynamical friction would be responsible for the lack of $L^\ast$ galaxies which is reflected in the requested magnitude gap of the observational definition. Moreover, to enhance the high efficiency in the merging process, FGs should have particular dynamical properties, such as the location of massive satellites on orbits with low angular momentum \citep[see][]{sommer06}. So, a combination of high mass satellite and low angular momentum orbits boosts the efficiency of the merging process \citep{boylan08}. Recently, \citet{lidman13} demonstrate that the growth of the BCGs since z$\sim$1 is mainly due to major mergers, suggesting that this could be the dominant mechanism in accreting the mass of central galaxies in cluster and thus supporting indirectly the merging scenario for fossil systems, which would differ from regular clusters only for the early time formation. Nevertheless, this evolutionary picture in which fossil groups became fossils in the early Universe and then evolved undisturbed is not the only proposed scenario. In the framework of the merging scenario \citet{diaz08} suggest that first ranked galaxies in fossil systems has the last major merger later than non-fossil ones. This means that the formation of large magnitude gaps as those in nowadays fossil systems is a long term process. In addition, \citet{vonbenda08} predict that the fossilness could be a transitional status of some systems. Thus, some fossil systems have become non-fossil ones in recent epoch due to accretion of nearby galaxy groups. An alternative formation scenario in which the magnitude gap of the systems appears at the beginning of the formation process can also explain the reported observational properties. This is the so called monolithic scenario, in which fossil systems formed with a top heavy LF. In this scenario, the magnitude gap is due to a primordial deficient formation of $L^\ast$ galaxies \citep{mul99}. All the observational results presented in the literature are limited by the small number of FGs known in the literature. A more general study of fossil systems is needed in order to discriminate between these two formation scenarios. For this reason, we started an extensive observational program called Fossil Groups Origins (FOGO), aimed at carrying out a large, systematic and multiwavelenght study of a sample of 34 FGs candidates identified by \citet{santos07}. The specific goals of the program include mass and dynamics of FGs, properties of their galaxy populations, formation of the central galaxies and their connection with the intragroup medium, properties of the extended diffuse light, and agreement with old and new simulations. The details of the project are resumed in the first paper of the series \citep[][]{aguerri11}. The structural properties of the BGGs in fossil and non-fossil systems were shown in the second paper \citep[][]{jairo12}. The L$_{\rm X}$-L$_{{\rm opt}}$ relation of fossil and normal systems will be presented in a forthcoming paper (Girardi et al., {\it in prep}). This is the third paper of the series, devoted to the characterization of the sample. In particular, we recomputed the magnitude gaps of the systems by using new spectroscopic redshift measurements. These new data provide us robust cluster membership and global properties for the systems. Only 15 out of 34 turned out to be fossil systems according with the standard definition \citep[see][]{jones03}. We have also analised the relations between central galaxies in FGs and non-FGs and other quantities such as magnitude gaps and velocity dispersion of the host halo. FGs follow the same relations than non-FGs, but they are extreme cases. The paper is organized as follows. The description of the sample is given in Section 2. The available dataset is shown in Section 3. Radial velocities determination are presented in Section 4. The results are given in Section 5. Sections 6 and 7 report the discussion and conclusions of the paper, respectively. Unless otherwise stated, we give errors at the 68\% confidence level. Throughout this paper, we use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}=0.7$ and $\Omega_M=0.3$ %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%% SECTION 2 - SAMPLE %% %%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | We characterized the sample of 34 FG candidates proposed by \citet{santos07} by using a unique collection of new optical photometric and spectroscopic data. This dataset was completed with SDSS-DR7 archival data. This large collection of radial velocities provided us robust cluster membership and global cluster properties for a subsample of 25 systems which were not available before. The fossilness determination of the 34 FG candidates was revisited. In particular, the magnitudes of the galaxies in each system were obtained by averaging three different magnitudes: Petrosian and model magnitudes from SDSS-DR7 and MAG-BEST SExtractor magnitude from our data. This was done because the magnitude of the BGGs can be affected by close satellites both in the SDSS and SExtractor analyses. Therefore, we computed new magnitude gaps ($\Delta m_{12}$ and $\Delta m_{14}$) within 0.5 $R_{200}$ for each system. The systems with $\Delta m_{12} \ge 2$ or $\Delta m_{14} \ge 2.5$ mag within the errors were classified as fossils. By applying this criterion, the total number of fossil systems in the sample is $15^{+8}_{-4}$. The uncertainties in the total number of fossil systems reflect the uncertainties in the $R_{200}$ determination. Moreover, there are 12 systems for which one or both the magnitude gaps are lower limits. For these systems, a more extended spectroscopic survey is needed in order to define their fossilness. We derived the main observational properties of the fossil systems in our sample. The fossil systems span a wide range of masses and we can confirm the existence of genuine fossil clusters in our sample. In particular, five fossil systems have LOS velocity dispersions $\sigma _{{\rm v}} > 700$ km s$^{-1}$, from both the $L_X$ luminosity and "shifting gapper" procedure. Clear correlations were found between the magnitude gaps and luminosity of the BGGs. In particular, the systems with larger $\Delta m_{12}$ have brighter BGGs, and the systems with larger $\Delta m_{12}$ or $\Delta m_{14}$ have larger fraction of the total galaxy light in the BGGs. The fossil systems also follow the same $L_{{\rm BGG}}/L_{{\rm tot}}-\sigma_{{\rm v}}$ relation of non-fossil systems. Nevertheless, they are extreme cases in the studied relations. In particular, the fossil systems have brighter BGGs than normal systems for any given LOS velocity dispersion (mass). All these properties can be explained by the two mainly accepted proposed scenarios of formation of fossil structures and thus are not conclusive in this sense. Nevertheless, we suggest that fossil systems with very bright central galaxies are not transitional phases of regular clusters and groups because, if this was the case, we should find systems with small gaps but very bright and massive central galaxies. These systems are not observed because the probability that two systems with such a bright BGG would merge is negligible. On the contrary, the systems with fainter BGGs possibly experienced a transitional fossil stage, which ended with the merging of another galaxy system. This could be the case of FGS06. The FOGO project will continue in the next future by analyzing other observational properties of fossil systems. In a forthcoming paper we will focus on the LFs of fossil and normal systems. This analysis will be crucial for the understanding of the formation and evolution of these galaxy aggregations, because the LFs of fossil systems in the merging scenario are expected to have a lack of $L^\ast$ galaxies. In contrast, the failed group formation scenario expects to find differences between fossil and normal systems in both the bright and faint ends of the LFs. | 14 | 3 | 1403.0588 |
1403 | 1403.4734_arXiv.txt | We present the results of sulfur monoxide, SO line emission observations of G0.253+0.016 with the Atacama Large Millimeter/submillimeter Array (ALMA) at an angular resolution of 1\arcsec.7. The dense and massive molecular cloud of G0.253+0.016 is highly sub-structured, yet shows no obvious signs of cluster formation. We found three outstanding features of the cloud from the SO emission, namely, shell structure of radius 1.3 pc, large velocity gradients of 20~km~s$^{-1}$~pc$^{-1}$ with the cloud, and cores with large velocity dispersions (30--40~km~s$^{-1}$) around the shell structure. We suggest that these large-velocity dispersion cores will form high-mass stars in the future. In attempt to explore the formation scenario of the dense cloud, we compared our results with numerical simulations, thus, we propose that G0.253+0.016 may have formed due to a cloud-cloud collision process. | Most stars, particularly high-mass stars ($>$$\,$8$~\MO$) form in stellar clusters (Lada $\&$ Lada 2003). Stellar clusters form in dense and massive molecular clumps (size$\sim$$\,$1~pc, mass$\sim$$\,$100--1000~$\MO$, density $\sim$10$^{3-5}$~$\rm{cm}^{-3}$) (Ridge et al. 2003; Lada $\&$ Lada 2003; Higuchi et al. 2009; Higuchi et al. 2010; Higuchi et al. 2013). Dense gas in the cluster forming clumps are significantly dispersed by the feedback from newly formed stars, thus blurring our understanding of cluster formation. To explore the initial conditions and details of cluster formation, it is necessary to study molecular clouds in the early stages of cluster formation. One of the unresolved issues of cluster formation is how dense clouds evolve into the compact and massive stellar clusters like the ``Arches" cluster (Lis $\&$ Menten 1998). The formation mechanism of such a super-star cluster is mysterious because they are expected to form in massive and small clouds of 1~pc scale. Therefore, without an external triggering process, it may be impossible to form super-star cluster. In recent literatures, cloud-cloud collisions are argued to be responsible for super-star cluster formation (Furukawa et al. 2009, Fukui et al. 2009, Ohama et al. 2010). Furukawa et al. (2009) suggested that young cluster, Westerlund 2, is likely to form by cloud-cloud collision mechanism. Westerlund 2 is a remarkable galactic super-star cluster, containing more than 10 massive stars with total stellar mass of 4500 $\MO$ in a small volume of only 1~pc in radius. The age of Westerlund 2 is suggested to be few Myrs (Ascenso et al. 2007; Furukawa et al. 2009; Rauw et al. 2007). G0.253+0.016 (hereafter G0.25) is very massive ($\sim$2$\times$10$^{5}\MO$, radius of 2.8~pc; Longmore et al. 2012) and dense ($\sim$3$\times$10$^{4}$~cm$^{-3}$; Kauffmann et al. 2013), with low dust temperature of 23~K (Rodr{\'{\i}}guez $\&$ Zapata 2013) and no obvious signs of cluster formation. G0.25 forms part of 100 pc circum-nuclear ring of clouds (Molinari et al. 2011) at 8.5$\,$kpc distance (Longmore et al. 2012). G0.25 presents an interesting site for the study of the initial condition of super-star cluster formation, because it is more massive and dense than the Orion A cloud, but hardly forms stars at all (Lis et al. 1994). The infrared luminosity of the entire cloud is 3$\times$10$^{5}$$\LO$, suggesting the presence of at least 5 embedded stars at evolutionary stages earlier than B0 (Lis et al. 2001). Rodr{\'{\i}}guez $\&$ Zapata (2013) also suggested that there are no O stars associated with the cloud. Recently, Kauffmann et al. (2013) presented N$_{2}$H$^{+}$ results arguing that G0.25 is presently far from forming high-mass stars and clusters. In this paper, we present 1\farcs7 resolution image of SO(v=0 3(2)--2(1)) line emission from the ALMA Cycle 0 observations. SO molecular line emissions are known to trace relatively dense and active clouds in a region (Bachiller et al. 2001; Aladro et al. 2013; Hacer et al. 2013; Codella et al. 2014). We investigated the spatial distributions and velocity structures with high-resolution observations based on ALMA archive data of G0.25 obtained with an extended configuration. The mosaic of the molecular line emission across this cloud were obtained at 90 GHz. We speculate that G0.25 cloud may represent the precursor to a super-star cluster. | We present the results of sulfur monoxide, SO, line emission observations of G0.25 with the ALMA, at an angular resolution of 1\arcsec.7. Our results and conclusions are summarized as follows: \begin{enumerate} \item We presented the SO map with a size of $\sim$ 3$^{\prime}$$\times$1.5$^{\prime}$ for G0.25 cloud. We identified detailed filamentary structures and cores within the cloud with high resolution observations. \item We discovered the shell structure of radius 1.3 pc with the SO emission. SO emission have been detected in some shocked regions, e.g., outflow shocked region in L\,1157 (Bachiller et al. 2001), tracing the cavity opened by the jet in NGC\,1333 (Codella et al. 2014), and NGC\,1068 (Aladro et al. 2013), implying that observed shell structure in G0.25 may have been due to shock related activity in the region. \item We found large velocity gradients of $\sim$ 20~km~s$^{-1}$~pc$^{-1}$ within the cloud, and cores with large velocity dispersions ($\sim$ 30 -- 40~km~s$^{-1}$) around the shell structure. We suggest that these high-velocity dispersion cores will form high-mass stars in the future. \item We estimated the virial ratios taking into account the contribution of rotation, and found that G0.25 cloud is likely to be gravitationally bound condition. However large velocity gradient in the G0.25 cloud is strange compared with the previous studies of the star forming regions. In order to explain the physical conditions of G0.25, we compared our results with numerical simulations, we propose that G0.253+0.016 may have formed due to a cloud-cloud collision process. \end{enumerate} \bigskip We thank the referee, Steven Stahler for constructive comments that helped to improve this manuscript. We also thank the ALMA staff for these observations during the commissioning stage. This paper makes use of the following ALMA data: ADS/JAO.ALMA 2011.0.00217.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. Finally, we acknowledge Elizabeth Tasker and Norikazu Mizuno for their contributions to our study. | 14 | 3 | 1403.4734 |
1403 | 1403.3578_arXiv.txt | The structure of the diffusion regions in antiparallel magnetic reconnection is investigated by means of a theory and a Vlasov simulation. The magnetic diffusion is considered as relaxation to the frozen-in state, which depends on a reference velocity field. A field-aligned component of the frozen-in condition is proposed to evaluate a diffusion-like process. Diffusion signatures with respect to ion and electron bulk flows indicate the ion and electron diffusion regions near the reconnection site. The electron diffusion region resembles the energy dissipation region. These results are favorable to a previous expectation that an electron-scale dissipation region is surrounded by an ion-scale Hall-physics region. [http://dx.doi.org/10.1063/1.4869717] | 14 | 3 | 1403.3578 |
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1403 | 1403.4672_arXiv.txt | We study the temporal and energy spectral properties of the unique neutron star low-mass X-ray binary \object{XTE J1701-462}. After assuming the HB/NB vertex as a reference position of accretion rate, the horizontal branch oscillation (HBO) of the HB/NB vertex is roughly 50 Hz. It indicates that the HBO is independent with the accretion rate or the source intensity. The spectral analysis shows $R_{\rm{in}}\propto\dot{M}_{\rm{Disk}}^{2.9\pm0.09}$ in the HB/NB vertex and $R_{\rm{in}}\propto\dot{M}_{\rm{Disk}}^{1.7\pm0.06}$ in the NB/FB vertex, which implies that different accretion rate may be produced in the HB/NB vertex and the NB/FB vertex. The Comptonization component could be fitted by constrained broken power law (CBPL) or nthComp. Different with \object{GX 17+2}, the frequencies of HBO positively correlate with the inner disk radius, which contradict with the prediction of Lense-Thirring precession model. \object{XTE J1701-462}, both in the Cyg-like phase and in the Sco-like phase, follows a positive correlation between the break frequency of broad band noise and the characteristic frequency of HBO, which is called the W-K relation. An anticorrelation between the frequency of HBO and photon energy is observed. Moreover, the rms of HBO increases with photon energy till $\sim \rm 10 ~\rm keV$. We discuss the possible origin of HBO from corona in \object{XTE J1701-462}. | Low mass X-ray Binary (LMXB) is composed of a compact object (neutron star or black hole) accreting matter from its low mass companion star ($\lesssim1M_{\odot}$). According to the X-ray spectral proporties and rapid timing variabilities, the neutron star (NS) LMXBs are usually classified as the Z sources and the atoll sources. They are named after the shapes which display in X-ray color-color diagram (CD) or hardness-intensity diagrams (HID). The Z sources produce approximate Eddington luminosities with soft X-ray spectra, whereas atoll sources produce a lower luminosities in the range $\sim 0.001-0.5 L_{\rm Edd}$ \citet{}. A typical Z source track shows three branches, from top to bottom, which are usually called the horizontal branch (HB), the normal branch (NB), and the flaring branch (FB; \citet{hasi89}), respectively. For atoll sources, the three branches are called extreme island, island, and banana state. Based on the orientation of branches, six typical Z sources are further divided into the Cyg- like Z sources (\object{Cyg X-2, GX 5-1}, and \object{GX 340+0}) with a horizontal HB (``Z"-shaped tracks) and the Sco-like Z sources (\object{Sco X-1}, \object{GX 17+2}, and \object{GX 349+2}) with a vertical HB (``$\nu$"-shaped tracks). The black hole (BH) and NS LMXBs show many similarities in their timing behaviors. Low-frequency quasi-periodic oscillations (LF-QPOs) are observed in Z sources, atoll sources and BH LMXBs. In general, the centroid frequencies of LF-QPOs are $\sim$1-70 Hz. The type C, B and A LF-QPOs in BH LMXBs were considered as corresponding to HBOs, NBOs, FBOs of Z sources \citep{cas05}, respectively. Moreover, the W-K relation, the strong correlation between the centroid frequency of LF-QPO and the break frequency in power density spectral, was identified in BH LMXBs, accreting millisecond pulsars and atoll sources \citep{wij99,bel02,str05}. Z sources show a similar but slightly shifted relation. These similar characteristics suggest that LF-QPOs are likely produced from the same physical mechanism in LMXBs. Lense-Thirring precession was introduced to interpret HBOs in NS LMXBs as well as type C LF-QPOs in BH LMXBs \citep{ste98,ste99}. In this model, the LF-QPOs were arisen from the misalignment between the compact star's spin axis and the rotational axis of the inner accretion disk \citep{len18,bar75}. \citet{ing09} discussed the possible origin of HBO from hot inner flow precession. The evolution of accretion rate $\dot{m}$ is studied from the HID of LMXB because it contains the variation of X-ray spectrum and radiation intensity. In BH LMXBs, $\dot{m}$ increases in the following direction: the Low Hard State (LHS) -- the Hard Intermediate State (HIMS) -- the Soft Intermediate State (SIMS) -- the High-Soft State (HSS) \citep{kal06,bel06,bel11}. Meanwhile, the type C LF-QPO was only observed in the HIMS in BH LMXBs. The frequency of type C QPO positively correlates with accretion rate and energy flux \citep{bel05}. However, the evolution of accretion rate $\dot{m}$ in the NS LMXBs is still controversial in the Cyg-like and Sco-like Z sources. The ascending trend of accretion rate is not always consistent with the source intensity increasing direction in the HIDs. According to multi-wavelength campaigns of the classic Z source Cyg X-2, accretion rate monotonically increases as HB-NB-FB \citep{has90,vrt90}. Based on the boundary layer emission model, \citet{gil05} also found $\dot{m}$ increasing from the HB to the FB in GX 340+0. However, Church and co-workers \citep{chu06, chu10} applied the extend ADC model for the Cyg-like Z sources and suggested an opposite direction, i.e., $\dot{m}$ increases from the FB/NB vertex to the HB/NB vertex. \citet{hom02} considered that $\dot{m}$ maintains constant along Z tracks. For classical Z sources, the frequency of HBO increased from HB to HB/NB vertex. So, three distinct correlations between the accretion rate and the frequency of HBO were proposed, that is, positive correlation, anticorrelation and non-correlation. The referred sources in the above works were analyzed either in the Cyg-like source or in the Sco-like source. The unique Z source, \object{XTE J1701-462}, switched from a Cyg-like Z source to a Sco-like Z source at high luminosity and from Z source to atoll source at low luminosity, which was observed by the Rossi X-ray Timing Explorer (RXTE) during its 2006-2007 outburst. The secular change of \object{XTE J1701-462} was driven by the accretion rate variation. \citet{lin09} studied the spectra evolutions of \object{XTE J1701-462} prudently, and suggested that the accretion rate $\dot{m}$ maintains constant on the NB and FB of Z sources. While on the HB, the ${L_{\rm{MCD}}-T_{\rm{MCD}}^{4/3}}$ correlation biased from the constant $\dot{m}$ line because the disk component encountered a Comptonization upscattering. The constant $\dot{m}$ should be satisfied after the Comptonization component accounted. In \object{GX 17+2}, the constant $\dot{m}$ was also established \citep{lin12}. \citet{hom10} indicated that the accretion rate was invariant in the Z sources and the oscillation of accretion rate $\dot{m}$ produced the Z tracks. However, \citet{ding11} concluded that the accretion rate of the disk follow $\dot{M}_{\rm{Disk}} \propto R^{7/2}_{\rm{in}}$ after considering the magnetic field effect during the accretion. In previous works, the relation between the spectra parameters and the characteristics of timing variability was not utilized to study the accretion rate variation. \object{XTE J1701-462} provides us a great opportunity to understand the temporal variabilities varying with the accretion rate evolution while the NS LMXB source transited from the Cyg-like, via the Sco-like Z source, to an atoll source. \citet{lin12} indicated two exactly opposite disk radius-HBO frequency relations in \object{GX 17+2} when the cutoff power law was replaced by nthComp. In this paper, we will provide a model independent method to study the HBOs' behaviors with decreasing accretion rate. In Sec. 2, we analyze the public archive data of \object{XTE J1701-462}. In Sec. 3, we study the X-ray spectra and timing variabilities of the HB/NB vertices and the NB/FB vertices, and then investigate the energy dependence of the HBO. The discussions and conclusions are displayed in Sec. 4 \& Sec. 5, respectively. | We study the HBOs of \object{XTE J1701-462} with decreasing accretion rate in the HB/NB vertex. We conclude that the HBO in \object{XTE J1701-462}, unlike the type C LF-QPOs in BH LMXBs, is independent with accretion rate. In other word, the ascending of HBO is not representing the accretion rate increasing. We also find that the anti-correlation relation between the frequency of HBO and its centroid energy. The energy dependence of HBO implies that the higher QPO produced a disk moves away from the NS. Both the Cyg-like phase and the Sco-like phase follow the W-K relation which are presumably caused by the same mechanism. The derived $R_{\rm in}-\nu_{\rm HBO}$ relations contradict the prediction of Lense-Thirring precession. We conclude that the HBO may origin from the corona. | 14 | 3 | 1403.4672 |
1403 | 1403.4358_arXiv.txt | Not only source catalogs are extracted from astronomy observations. Their sky coverage is always carefully recorded and used in statistical analyses, such as correlation and luminosity function studies. Here we present a novel method for catalog matching, which inherently builds on the coverage information for better performance and completeness. A modified version of the Zones Algorithm is introduced for matching partially overlapping observations, where irrelevant parts of the data are excluded up front for efficiency. Our design enables searches to focus on specific areas on the sky to further speed up the process. Another important advantage of the new method over traditional techniques is its ability to quickly detect dropouts, i.e., the missing components that are in the observed regions of the celestial sphere but did not reach the detection limit in some observations. These often provide invaluable insight into the spectral energy distribution of the matched sources but rarely available in traditional associations. | \label{sec:intro} \noindent With the technology improvements in modern telescopes over the last decade, astronomy now has large volumes of data that will soon reach the peta-byte regime. Novel data storage, automated processing and sophisticated publishing solutions are being and will have to be developed to handle the challenges of ongoing and next-generation surveys, such as the Large Synoptic Survey Telescope \citep[LSST;][]{2006SPIE.6270E..24B}. Scientific progress is increasingly dependent upon our ability to analyze the immense amount of measurements to extract more information and ultimately develop an understanding of Nature's fundamental properties. Today many specialized instruments and facilities release invaluable data to the public on a regular basis. Multicolor and time-series studies rely on all of these observations. One of the critical requirements of these analyses is automated tools that can meaningfully combine the data sets from several independent archives \citep{2001SPIE.4477...20M}. In particular, catalog matching has been the focus of several studies. Some of these revisit the statistical aspects of finding the right discriminators for the associations \citep{2008ApJ...679..301B}, while others focus on the computational issues \citep{2005ApJ...622..759G, 2007cs........1171G, 2011ASPC..442...85P} to beat down the (naively) combinatorial nature of the problem. These efforts are fundamentally influencing the way we handle observations and pave the road for next-generation analysis tools and services, such as those of the virtual observatories. Folding in the sky coverage into catalog matching is key to efficient algorithms. Previous approaches either completely neglect to use this information or include them as an afterthought. By ignoring the footprints we eliminate our ability to constrain the brightness of non-detections that might appear in other observations. Without the coverage information, we simply cannot tell whether a missing match means a lower flux than the observational limit or the direction is simply outside the observed area. There are spatial query engines \citep[e.g.,][]{2001misk.conf..631K, 2004ASPC..314..289F} that can quickly find sources in a given region of the sky but they typically use a different indexing scheme to speed up the searches, and therefore require extra steps to re-organize the data set for the cross-matching. Our approach is different. We create a procedure that incorporates spatial constraints directly into the crossmatching. It uses the same indexing scheme, hence performs better than previous techniques and can also deliver dropouts for further constraints on the spectral energy distribution of the sources. In this paper, we introduce a solution that builds on the sky coverage information of the catalogs to exclude irrelevant parts of the data sets before the matching begins. The method is introduced in Section~\ref{sec:method}. Section~\ref{sec:implement} describes the database implementation of the method, and extended to efficiently detect dropouts, Section~\ref{sec:perform} discusses the performance, and Section~\ref{sec:con} concludes our study. | \label{sec:con} \noindent We introduced a new algorithm for catalog matching that inherently uses the sky coverage information of all observations. The key feature of our approach is that it incorporates the coverage straight into the crossmatching and performs the filtering at runtime. It automatically restricts the datasets to the intersecting area to gain performance. User-defined spatial constraints can also be input for targeted matching. The developed method builds on the Zones Algorithm. The footprints of the catalogs are approximated as R.A.\ intervals within the thin zones. The representation is not only highly accurate but also in line with the data organization required for efficient matching. Our implementation is in SQL, the Structure Query Language that is understood by most astronomy archives today including the SDSS, GALEX and HLA. Coding details are also provided along with measures of the performance scaling as a function of sky coverage. The method also enables us to efficiently detect missing observations that were covered but did not appear over the detection threshold. It is a fast though preliminary step to identify dropouts. These dropouts are critical to constrain the spectral energy distribution of the celestial objects. Knowing that the flux is below a given limit can be very informative for a variety of studies ranging from photometric redshift estimation to galaxy evolution and cosmological analyses. | 14 | 3 | 1403.4358 |
1403 | 1403.1261_arXiv.txt | We study cosmic metal enrichment via AMR hydrodynamical simulations in a $(10$~Mpc~$h^{-1})^{3}$ volume following the Pop III -- Pop II transition and for different Pop III IMFs. We have analyzed the joint evolution of metal enrichment on galactic and intergalactic scales at $z=6$ and $z=4$. Galaxies account for $\lsim9\%$ of the baryonic mass; the remaining gas resides in the diffuse phases: (a)~\textit{voids}, i.e. regions with extremely low density ($\Delta\leq 1$), (b) the true \textit{intergalactic medium} (IGM, $1<\Delta\leq 10$) and (c) the \textit{circumgalactic medium} (CGM, $10<\Delta\leq 10^{2.5}$), the interface between the IGM and galaxies. By $z=6$ a galactic mass-metallicity relation is established. At $z=4$, galaxies with a stellar mass $M_{\star}\simeq10^{8.5} M_\odot$ show $\log({\rm O}\slash H)+12=8.19$, consistent with observations. The total amount of heavy elements rises from $\Omega^{\rm SFH}_{Z}=1.52\times10^{-6}$ at $z=6$ to $8.05 \times10^{-6}$ at $z=4$. Metals in galaxies make up to $\simeq0.89$ of such budget at $z=6$; this fraction increases to $\simeq0.95$ at $z=4$. At $z=6$ ($z=4$) the remaining metals are distributed in CGM/IGM/voids with the following mass fractions: $0.06/0.04/0.01$ ($0.03/0.02/0.01$). Analogously to galaxies, at $z=4$ a density-metallicity ($\Delta-Z$) relation is in place for the diffuse phases: the IGM/voids have a spatially uniform metallicity, $Z\sim10^{-3.5}\zsun$; in the CGM $Z$ steeply rises with density up to~$\simeq10^{-2}\zsun$. In all diffuse phases a considerable fraction of metals is in a warm/hot ($T\, \mu^{-1}>10^{4.5}K$) state. Due to these physical conditions, $\CIV$ absorption line experiments can probe only $\simeq$ 2\% of the total carbon present in the IGM/CGM; however, metal absorption line spectra are very effective tools to study reionization. Finally, the Pop III star formation history is almost insensitive to the chosen Pop III IMF. Pop III stars are preferentially formed in truly pristine ($Z=0$) gas pockets, well outside polluted regions created by previous star formation episodes. | The intergalactic medium (IGM) has been extensively investigated through the study of the $\HI$ Ly$\alpha$ forest \citep{Rauch:1998} and the absorption features due to ionized metal species \citep[e.g.][]{Songaila:1996} detected in the spectra of high redshift quasars (QSO). Observations have probed metal enrichment in different intergalactic environments: damped Ly$\alpha$ absorbers (DLA), characterized by column densities $\log N_{\rm HI}\slash{\rm cm}^{-2}\gsim 20$ and metallicities $10^{-1.5}\lsim Z\slash\zsun\lsim10^{-1}$; Lyman limit ($17\lsim\log N_{\rm HI}\slash{\rm cm}^{-2}\lsim 20$) and Ly$\alpha$ forest ($14\lsim\log N_{\rm HI}\slash{\rm cm}^{-2}\lsim 17$) systems, typically enriched at $10^{-3.5}\lsim Z\slash\zsun\lsim10^{-2}$ \citep{Meiksin:2009RvMP}. The evolution of the IGM enrichment can be studied by measuring the abundance of ionized species at different cosmic times. For example, the $\CIV$ density parameter decreases with redshift from $\Omega_{\rm CIV}\simeq 8\times10^{-8}$ at $z\simeq 0$ to $\Omega_{\rm CIV}\simeq 10^{-8}$ at $z\simeq 2.5$ \citep[e.g.][]{DOdorico:2010}, it is constant up to $\lsim5$ \citep[e.g.][]{Schaye:2003ApJ,Cooksey:2010ApJ}, and could possibly show sign of a downturn, $\Omega_{\rm CIV}\lsim0.5\times10^{-8}$, for $z\gsim 5$ \citep[e.g.][]{Ryan-Weber:2009MNRAS,Becker:2009ApJ,Simcoe:2011ApJ,DOdorico:2013MNRAS}. Similarly, $\SIV$ displays a flat behavior for $2.5\lsim z \lsim5$ \citep{Songaila:2001ApJ,Songaila:2005AJ}. Metals are produced by stars inside galaxies featuring a cosmic star formation rate density (SFR) $\simeq 10^{-2} \msun {\rm yr}^{-1} {\rm Mpc}^{-3}$ at $z=0$, increasing up to one order of magnitude at $z\sim 3$ and decreasing by a factor $\sim10^{2}$ for $3\lsim z \lsim9$ \citep{Dunlop:2013ASSL}. IGM is metal polluted by galaxies, whose SFR can be can used to infer the total metal density parameter. \citet{Pettini:1999cezh} have shown that observations account for $\sim20\%$ of metals at $z\gsim2$, implying that the missing ones must be locked in a warm-hot ionized phase \citep[e.g.][]{Ferrara:2005ApJ}. Galaxies with high ($M_{\star}\gsim10^{11}\msun$) and low ($M_{\star}\lsim10^{8}\msun$) stellar mass display different characteristics. The local \citep{Panter:2008MNRAS} and $z\lsim3$ \citep{Maiolino:2008A&A,Mannucci:2010MNRAS} mass-metallicity ($\massmetal$) relation shows an increasing metal abundance with $M_{\star}$ from $\sim10^{4}\msun$ \citep{Kirby:2013arXiv} up to $\sim10^{10}\msun$ and a flattening for higher stellar mass. Similar difference is present in the dark-to-visible mass ratio \citep[e.g.][]{Guo:2010MNRAS,Tolstoy:2010IAUS,McGaugh:2010ApJ}: dwarf galaxies typically show values $\gsim30$, while Milky Way size galaxies have a ratio $\lsim15$; for increasing mass the value approaches (but does not reach) the cosmological one, $\Omega_{m}\slash\Omega_{b}\sim 6$. A theoretical framework must auto-consistently account for the history of IGM enrichment, its thermal state and the global evolution of galaxy formation. An attractive scenario consists in the so-called {\it pre-enrichment} \citep{Madau:2001ApJ,Ferrara:2008IAUS}, in which IGM pollution is mainly due to low mass ($M_{\star}\lsim10^{7}\msun$) galaxies ejecting metals via supernova (SN) explosions at high redshift ($z\gsim8$). In this picture massive galaxies are able to retain their metals, thus their evolution follows a closed box chemical model\citep[e.g.][]{Tremonti:2004ApJ}. On the other hand, low mass galaxies are prone to material ejection because of the shallower potential well and since their smaller size allows the SN to coherently drive the outflows \citep[e.g.][]{Ferrara:2000MNRAS}. Ejection from low mass galaxies has an obvious advantage in terms of volume filling factor, given the abundance of these sources in $\Lambda$CDM models\footnote{In this work we assume a $\Lambda$CDM cosmology with total matter, vacuum and baryonic densities in units of the critical density $\Omega_{\Lambda}= 0.727$, $\Omega_{dm}= 0.228$, $\Omega_{b}= 0.045$, Hubble constant $\rm H_0=100~h~km~s^{-1}~Mpc^{-1}$ with $\rm h=0.704$, spectral index $n=0.967$, $\sigma_{8}=0.811$ \citep[][]{Larson:2011}.} \citep[i.e.][]{Press:1974}. Additionally, an early ($z\gsim5$) IGM pollution allows the shocked and enriched gas to cool down \citep[e.g.][]{Ferrara:2008IAUS}, which can explain the observed narrowness ($\sim15~{\rm km~s}^{-1}$) of the Doppler width of metal lines \citep{Meiksin:2009RvMP}. The {\it pre-enrichment} scenario is appealing because the same low mass galaxies which start to pollute the IGM can play an important role in the firts stages ($z\gsim8$) of cosmic reionization \citep[e.g.][]{Salvadori:2013arXiv}. In particular, \citep[e.g.][]{Choudhury:2008MNRAS} low mass galaxies are the ideal hosts for first metal-free Population III (PopIII) stars, which may possibly be the responsible for an early ($z\simeq7$) reionization \citep[e.g.][]{Gallerani:2008MNRASa,Gallerani:2008MNRASb}. However, the nature of Pop III stars is still under debate, because of lack of observations. From a theoretical point of view there is no clear consensus on their formation properties \citep{Bromm:2002ApJ,Yoshida:2006ApJ,Greif:2012MNRAS,Hosokawa:2012ApJ,Meece:2013arXiv} nor their subsequent evolution \citep{Heger:2002ApJ,Nomoto:2006NuPhA}. Thus, understanding the IGM metal enrichment from the first galaxies is fundamental to explain both the formation and the evolution of galaxies and the reionization history \citep{Barkana:2001PhR,Ciardi:2005SSR}. However, it remains to be assessed the evolution of the temperature and chemical state of the enriched IGM, the imprint the metal transport has on the formation of new star forming regions and its role in the transition from Pop III to the successive generation of Population II (PopII) stars. The steady increase of availability of high redshift IGM observations \citep[e.g.][]{DOdorico:2013MNRAS} and their interpretation can hopefully clarify the picture, by better constraining the theoretical models. Cosmological numerical simulations have been extensively used to study the problem \citep{Aguirre:2007EAS,Johnson:2011arXiv}. However, the huge dynamical range of the underlying physical phenomena makes a true auto-consistent simulation impossible. A viable modelization can be achieved by using subgrid models. These depend both on the considered physics and code implementation. Recently, \citet{Hopkins:2013arXiv} studied the impact of different star formation criteria, \citet{Agertz:2012arXiv} and \citet{Vogelsberger:2013arXiv} analyzed the effect of including different kind of feedback, and the AQUILA project \citep[i.e.][]{aquila:2012MNRAS} compared 13 different prescriptions of the main used cosmological codes. Subgrid modelling lessens the burden of the large dynamical range, but given the currently available computational capabilities the numerical resources have to be focused toward either the small or the large scales. Simulations of small cosmic volumes, i.e. box sizes $\lsim\, {\rm Mpc}\,h^{-1}$, concentrate the computational power and allow the usage of highly refined physical models. \citet{Greif:2010ApJ} studied the transition from Pop III to Pop II stars in a $10^{8}\msun$ galaxy at $z\sim10$ assessing the role of radiative feedback; \citet{Maio:2010MNRAS} analyzed the same transition by varying several parameters, such as the critical metallicity $\zcrit$ that distinguishes the populations, the initial mass function (IMF), the metal yields and the star formation threshold; \citet{xu:2013arXiv} focalized on pinpointing the remnant of Pop III at high redshift, by employing the same computational scheme of \citet{Wise:2012ApJ}, which analyzed the impact of radiation from first stars on metal enrichment at $z\gsim9$; at the same redshift, \citet{Biffi:2013arXiv}, using an extensive chemical network, studied the properties and the formation of first proto-galaxies. Large scale ($\gsim5\, {\rm Mpc}\,h^{-1}$) cosmological simulations naturally allows for a fair comparison with the observations. \citet{Scannapieco:2006MNRAS} showed that observation of line of sight (l.o.s.) correlations of $\CIV$ and $\SIV$ are consistent with a patchy IGM enrichment, confined in metal bubbles of $\sim 2\,{\rm Mpc}\,h^{-1}$ at $1.5\lsim z\lsim3$; by implementing galaxy outflows driven by a wind model \citet{Oppenheimer:2006MNRAS} managed to reproduce the flatness of $\Omega_{\rm CIV}$ at $2\lsim z\lsim5$; \citet{Tornatore:2007MNRAS} found evidence of Pop III production at $z\gsim4$, hinting at the possibility of observing metal-free stars; using a galactic super wind model \citet{Cen:2011ApJ} simulated a 50 Mpc~$h^{-1}$ box finding, among other results, a good agreement with observations for $\Omega_{\rm CIV}$ and a reasonable match for $\Omega_{\rm OVI}$; by using a $(37.5\,{\rm Mpc}\,h^{-1})^{3}$ volume simulation evolved up to $z=1.5$, and considering different IMFs and feedback mechanisms, \citet{Tescari:2011MNRAS} analyzed the evolution of $\Omega_{\rm CIV}$ and statistics of \HI\, and $\CIV$ absorbers at different redshifts; simulating a box with size of $25\,{\rm Mpc}\,h^{-1}$ and including various feedback, \citet{Vogelsberger:2013arXiv} managed to match several observations, as the SFR and stellar mass density (SMD) evolution for $z\lsim 9$, the galaxy stellar mass function and mass-metallicity relation at $z=0$. The aim of this paper is to model the IGM metal enrichment focusing on high redshift ($z\geq4$) by simulating a volume large enough to include a statistically significant ensemble of galaxies. Clearly, the trade off consists in a limitation of the resolution and small scale complexity that can be investigated. Our modelling approach is to limit the number of free parameters of the subgrid prescriptions and constrain them with first galaxies observations, namely both global SFR densities inferred from Ultraviolet (UV) luminosity functions \citep{Bouwens:2012ApJ,Zheng:2012Natur} and SMD from stellar energy density fitting \citep{Gonzalez:2011}. This method limits the uncertainty on the feedback prescriptions \citep[e.g.][]{Vogelsberger:2013arXiv}, and at the same time it allows a large scale analysis of the metal enrichment process. The paper is structured as follows. In Sec. \ref{sec_metodo_sim}, we describe the numerical implementation of the cosmological simulations whose free parameters are then calibrated by matching SFR and SMD data in Sec. \ref{sec_sfr_smd}. Sec. \ref{sec_gal_enrichment} and Sec. \ref{sec_global_result} contain the analysis of galactic and IGM metal enrichment, respectively. We devote Sec. \ref{sec_test} to study the effects of varying the Pop III IMF, and in Sec. \ref{sec_spettri} we compute and discuss mock QSO absorption spectra, in preparation for a future detailed comparison recent of high redshift ($4\lsim z\lsim6$) absorption line data \citep{DOdorico:2013MNRAS}. Finally, in Sec. \ref{sec_conclusioni} we present our conclusions. \begin{table} \centering \begin{tabular}{|c|c|c|c|} \hline $\log \left(Z\slash \zsun\right)$ & $Y$ & $R$ & $\epsilon_{\popii}$\\ \hline\hline $-4.0$ & 0.0160 & 0.4680 & $10^{50}$\\ $-2.0$ & 0.0192 & 0.4705 & $10^{50}$\\ $-1.0$ & 0.0197 & 0.4799 &$10^{50}$\\ $\,0.0$ & 0.0253 & 0.4983 &$10^{50}$\\ \hline \end{tabular} \caption{Adopted IMF-averaged Pop II metal yields and gas return fractions (\citet{vandenHoek:1997A&AS} for $0.8 \le m/\msun \le 8 $ and \citet{Woosley:1995ApJS} for $8 \le m/\msun \le 40 $); explosion energies (in \mbox{erg\,}$\msun^{-1}$) are taken from \citet{Woosley:1995ApJS}. \label{tabella_yield}} \end{table} \begin{table} \centering \begin{tabular}{|c|c|c|c|} \hline IMF & $Y$ & $R$ & $\epsilon_{\popiii}\slash\epsilon_{\popii}$ \\ \hline\hline SALP & 0.0105 & 0.46 & 1 \\ FHN & 0.0081 & 0.76 & 1 \\ PISN & 0.1830 & 0.45 & 10 \\ \hline \end{tabular} \caption{Adopted IMF-averaged Pop III metal yields and gas return fractions, and relative explosion energies. Data are taken from \citet{vandenHoek:1997A&AS} for $0.8 \le m/\msun \le 8 $ and \citet{Woosley:1995ApJS} for $8 \le m/\msun \le 40 $ for SALP, \citet{Kobayashi:2011ApJ} for FHN, and \citet{Heger:2002ApJ} for PISN. \label{tabella_yield_popiii}} \end{table} | We have studied cosmic metal enrichment via a suite of $\Lambda$CDM hydrodynamical simulations using a customized version of the adaptive mesh refinement code {\tt RAMSES} to evolve a $(10$~Mpc~$h^{-1})^{3}$ volume up to $z=4$ with $512^{3}$ dark matter (DM) particles, a corresponding number of coarse grid cells and allowing for 4 additional levels of refinement. The subgrid prescription for star formation is based on a local density threshold criterion ($\Delta>\Delta_{\rm th}$) and on a critical metallicity criterion ($\zcrit = 10^{-4}\zsun$), allowing us to follow the transition from Pop III to Pop II stars. To assess the impact of variations in the unknown Pop III IMF we have investigated three different choices: (a) a standard Larson-Salpeter IMF (SALP), (b) a $\delta$-function describing faint hypernovae (FHN), and (c) a top-heavy IMF allowing for pair-instability supernovae (PISN). We account for thermal feedback from supernovae and implemented a metal-dependent parameterization of stellar yields and return fractions. This set-up enables the resolution of DM halos masses of $10^{7.5}\msun$ with $\simeq 100$ particles and to build a statistically significant sample of galaxies at all redshifts of interest. The two free parameters of our subgrid model (star formation timescale and supernova coupling efficiency) have been fixed by reproducing the observed cosmic star formation rate \citep[SFR,][]{Bouwens:2012ApJ,Zheng:2012Natur} and stellar mass densities \citep[SMD,][]{Gonzalez:2011} at $4\leq z \lsim 10$. By constructing halo catalogues and identifying the associated stars and star forming regions ($\Delta>\Delta_{\rm th}$), it has been possible to analyze the evolution of metal enrichment on galactic scales at two representative redshifts, $z=6$ and $z=4$. Galaxies account for $\lsim 9\%$ of the baryonic mass; the complementary fraction resides in the diffuse medium, which we have classified according to the environmental overdensity into: (a)~\textit{voids}, i.e. regions with extremely low density ($\Delta\leq 1$), (b) the true \textit{intergalactic medium} (IGM, $1<\Delta\leq 10$) and (c) the \textit{circumgalactic medium} (CGM, $10<\Delta\leq 10^{2.5}$), representing the interface between the IGM and galaxies. We have computed synthetic spectra of metal absorption lines through the simulated box at $z=6$. The number density of different ionic species are calculated in post-processing with {\tt CLOUDY} and by considering two physically motivated and observationally constrained reionization models, i.e. an Early Reionization Model (ERM, $\log(\Gamma_{\rm HI}/{\rm s^{-1}})=-12.46$ at $z=6$) and a Late Reionization Model (LRM, $\log(\Gamma_{\rm HI}/{\rm s^{-1}})=-12.80$). We have tried to analyze separately the metal enrichment properties of galaxies and diffuse medium for sake of clarity, but obviously the intimate connection between these two components makes it impossible to separate their description completely. Readers mostly interested in galaxies/stars (diffuse gas) can directly refer to Sec. \ref{sec_gal_enrichment} (Sec. \ref{sec_global_result}); those specifically interested in Pop III stars should also find Sec. \ref{sec_test} relevant. The summary of the main results given below is organized in points attempting to keep these distinctions. \begin{itemize} \item[\bf 1.] Between $z=9$ and $z=6$ a galactic mass-metallicity relation is established. For star forming regions of mass $M_{\rm SF}\gsim 10^{7} \msun$, such relation shows little evolution from $z=6$ to $z=4$. In particular, at $z=4$, galaxies hosting a stellar mass $M_{\star}\simeq10^{8.5}\msun$ show a mean oxygen abundance of $\log({\rm O}\slash H)+12=8.19$, consistent with observations \citep[][]{Troncoso:2013arXiv1311}. \item[\bf 2.] At $z = 4$ such relation extends to $M_{\rm SF}\lsim 10^{7} \msun$: these are satellite galaxies forming whose star formation has been enabled by the progressive enrichment of the diffuse gas out of which they form. For $10^{6}\lsim M_{\rm SF}\slash\msun\lsim 10^{7}$ the metallicity trend is flat and resembles the one observed in the faintest Local Group dwarf galaxies \citep[e.g.][]{Kirby:2013arXiv}. \item[\bf 3.] The total amount of heavy elements produced by star formation rises from $\Omega^{\rm SFH}_{Z} = 1.52 \times 10^{-6}$ at $z=6$ to $8.05 \times 10^{-6}$ at $z=4$. Metals in galaxies make up to $\simeq 0.89$ of such budget at $z=6$; this fraction increases to $\simeq 0.95$ at $z=4$. At $z=6$ ($z=4$) the remaining metals are distributed in the three diffuse phases, CGM/IGM/voids, with the following mass fractions: $0.06/0.04/ 0.01$ ($0.03/0.02/ 0.01$). \item[\bf 4.] In all the diffuse phases a considerable fraction of metals is in a warm/hot ($T\, \mu^{-1}>10^{4.5}K$) state. In particular, a small but not negligible mass fraction ($\simeq0.003$) of metals in voids shows $T\, \mu^{-1}\leq10^{4.5}K$. This implies that these metals must have been injected at sufficiently early epochs that they had the time to cool as expected in a pre-enrichment scenario. \item[\bf 5.] Analogously to the mass-metallicity relation for star forming regions, at $z=4$ a density-metallicity ($\Delta-Z$) relation is in place for the diffuse phases. Independently of $\Delta$, the IGM/voids show an uniform distribution around $Z\sim10^{-3.5}\zsun$, while in the CGM $Z$ steeply rises with density up to~$\simeq10^{-2}\zsun$. \item[\bf 6.] The geometry of metal bubbles is influenced by the topology of the cosmic web. At $z=6$, $\sim40\%$ are spherically symmetric and are mostly found around isolated galaxies; $30\%$ show instead a cylindrical shape which mainly results from merging of bubbles aligned along filaments. \item[\bf 7.] The cosmic Pop III star formation history is almost insensitive to the chosen Pop III IMF. Pop III stars are preferentially formed in pockets of pristine ($Z=0$) gas, well outside polluted regions created by nearby/previous star formation episodes. This supports the ``Pop III wave'' scenario suggested by \citet{Tornatore:2007MNRAS} and confirmed by \citet{Maio:2010MNRAS}. \item[\bf 8.] In the PISN case, the Pop II SFR is suppressed by a factor of $\simeq5$ with respect to the SALP/FHN cases. Because of the higher energy deposition, a pair-instability SN can reach and disrupt a nearby potential star formation site, quenching Pop II formation. Assuming the same star formation timescales for Pop II and Pop III, a PISN scenario is difficult to be reconciled with the observed SFR history, as the feedback from these stars is probably too effective. \item[\bf 9.] Metal absorption line spectra extracted from our simulations at $z\sim 6$ contain a greater wealth of information with respect to the Ly$\alpha$ forest. Given the prevailing thermodynamical/ionization conditions of the enriched gas, $\CIV$ absorption line experiments can only probe up to $\simeq$ 2\% of the total carbon present in the IGM/CGM. However, metal absorption lines are very effective tools to study reionization. \item[\bf 10.] The occurrence of low-ionization metal systems (e.g. $\OI$ and $\CII$) in $z\sim 6$ quasar (gamma-ray burst) absorption spectra does not exclude the possibility that the IGM/CGM is on average highly ionized at these epochs. In fact, such systems, although with a lower incidence than in a Late Reionization Model, are also detectable in the Early Reionization Model, which predicts a lower $\HI$ fraction ($x_{\rm HI}\sim 10^{-4}$) at $z\simeq 6$. \end{itemize} In the future, we will perform a more extended statistical analysis of the synthetic spectra, in terms of the equivalent width and column density distributions. This study will enable a direct comparison with recent high-$z$ observations \citep{DOdorico:2013MNRAS} and will allow to constrain cosmic reionization models. Since the ionization level of metal atoms is sensitive to the proximity effect of ionizing sources, it will be crucial to take into account radiative transfer effects. Therefore, we plan to couple our simulation with the new version of the radiative transfer code {\tt CRASH3} \citep{Graziani:2013MNRAS}. This code allows to include an arbitrary number of point sources and to reprocess the ionizing radiation through an inhomogeneous distribution of metal-enriched gas, therefore representing a perfect tool for our planned research. As a final remark, we have highlighted the potential problem that chemical feedback might be artificially enhanced in a simulation when the box size becomes smaller or comparable to the pollution radius $\langle R_B \rangle$. Although the box size and resolution have a significant impact on the determination of Pop III cosmic SFR, additional uncertainties come from the treatment of radiative feedback. A proper demonstration would involve a suite of simulations with increasing box size and fixed resolution and a convergence study with fixed box-size and increasing mass resolution. The situation may be worth a closer scrutiny, which we defer to future work. | 14 | 3 | 1403.1261 |
1403 | 1403.1782_arXiv.txt | We consider axino warm dark matter in a supersymmetric axion model with R-parity violation. In this scenario, axino with the mass $m_\axino\simeq 7$ keV can decay into photon and neutrino resulting in the X-ray line signal at $3.5\kev$, which might be the origin of unidentified X-ray emissions from galaxy clusters and Andromeda galaxy detected by the XMM-Newton X-ray observatory. | \label{introduction} Various astrophysical and cosmological observations provide convincing evidences for the existence of dark matter (DM). Dark matter distribution spans in wide range of scales from galaxy to clusters of galaxies and the large scale structure of the Universe. Recently, anomalous X-ray line emissions have been observed from galaxy clusters and also in the Andromeda galaxy~\cite{Bulbul:2014sua,Boyarsky:2014jta}. While those might be a result of systematic effects, it would be interesting if the line came from the new source of astrophysical phenomena or from new physics. It was suggested that the signal might come from decaying dark matter with the mass and lifetime, \dis{ m_{\rm DM} &\simeq 7\kev,\\ \tau_{\rm DM} & \simeq 2\times 10^{27} - 2\times 10^{28} \, \sec, \label{xray} } assuming that they are the dominant component of dark matter. Some theoretically interesting particle models have been suggested such as sterile neutrino~\cite{Bulbul:2014sua,Boyarsky:2014jta,Ishida:2014dlp,jian:2014gza}, exciting dark matter~\cite{Finkbeiner:2014sja}, millicharged dark matter~\cite{Aisati:2014nda,Frandsen:2014lfa}, axion like particle~\cite{Higaki:2014zua,Jaeckel:2014qea,Lee:2014xua}, in the effective theory~\cite{Krall:2014dba}. In this paper, we study the warm dark matter axino in a R-parity violating supersymmetric model. With bilinear R-parity breakings, neutralinos mix with neutrinos and thus the axino can decay into photon and neutrino. We find that the axino mass with $7$ keV can have the proper lifetime and relic density for the X-ray line emission. In this scenario, as an interesting consequence, the upper bound on the neutrino mass imposes that the Bino mass is lighter than about $10$ GeV. In Section~\ref{axino} we introduce the model of axino dark matter and in Section~\ref{Rparity} we consider the R-parity violation and decay of axinos. We summarize in Section~\ref{conclusion}. | \label{conclusion} Axino is a good candidate for dark matter. When its mass is around $7\kev$ and R-parity is broken bilinearly, the axino decays into photon and neutrino. We studied this decaying axino warm dark matter in the light of the recent observation of X-ray line emission from the center of galaxy clusters and Andromeda galaxy observed by XMM-Newton. We find that the decaying axino can naturally explain the X-ray signal with/without additional component of dark matter. We note that the neutrino mass bound implies that the Bino mass is less than about $10$ GeV. | 14 | 3 | 1403.1782 |
1403 | 1403.7098_arXiv.txt | We simulate the adiabatic contraction of a dark matter (DM) distribution during the process of the star formation, paying particular attention to the phase space distribution of the DM particles after the contraction. Assuming the initial uniform density and Maxwellian distribution of DM velocities, we find that the number $n(r)$ of DM particles within the radius $r$ scales like $n(r) \propto r^{1.5}$, leading to the DM density profile $\rho\propto r^{-1.5}$, in agreement with the Liouville theorem and previous numerical studies. At the same time, the number of DM particles $\nu(r)$ with periastra smaller than $r$ is parametrically larger, $\nu(r) \propto r$, implying that many particles contributing at any given moment into the density $\rho(r)$ at small $r$ have very elongated orbits and spend most of their time at distances larger than $r$. This has implications for the capture of DM by stars in the process of their formation. As a concrete example we consider the case of primordial black holes (PBH). We show that accounting for very eccentric orbits boosts the amount of captured PBH by a factor of up to $2\times 10^3$ depending on the PBH mass, improving correspondingly the previously derived constraints on the PBH abundance. | \label{sec:introduction} Astrophysical and cosmological observations have provided a compelling evidence that about 28\% of the energy density of the Universe is contained in the form of a non-relativistic non-baryonic dark matter (DM) \cite{2013arXiv1303.5076P}. Despite the extensive experimental efforts, all the attempts at direct and indirect non-gravitational detection of this matter have been unsuccessful so far, so that the DM nature remains essentially unconstrained, leaving room for a diversity of candidates. It is often assumed that the DM is composed of some kind of new particles beyond the Standard Model --- axion-like particles, sterile neutrinos, weakly interacting massive particles are the most common examples. However, other candidates such as primordial black holes (PBH) may provide a viable alternative. In the latter case, the advantage is that no new particles beyond the Standard Model are required. Apart from direct and indirect searches, competitive constraints on DM properties may be obtained from observations of stars where the DM could have been accumulated and produced observable effects. In the Sun, the DM annihilation may result in the observable flux of high energy neutrinos \cite{2009JCAP...08..037N}. DM may also induce abnormal asteroseismological effects~\cite{Lopes:2012af} or suppress convection zones~\cite{Casanellas:2012jp,Casanellas:2013nra} or even modify the transportation properties throughout the star~\cite{Frandsen:2010yj,Horowitz:2012jd}. More catastrophic effects may result from the DM accumulation in compact stars such as neutron stars (NSs) or white dwarfs(WDs). If the particles of DM are not self-annihilating (e.g., asymmetric dark matter), the accumulated amount of DM may become sufficient to form a black hole (BH) inside the compact star~\cite{Kouvaris:2010jy,Kouvaris:2010vv,Kouvaris:2011fi}. Because of a much higher density of nuclear matter in compact stars as compared to main sequence stars, the accretion is sufficiently efficient to destroy the star in a short time (see Ref.~\cite{Kouvaris:2013kra} for a detailed discussion, including the role of the angular momentum). In this case mere observations of compact stars imply constraints on the DM properties~\cite{Kouvaris:2010jy,Kouvaris:2010vv,Kouvaris:2011fi}. The same considerations obviously apply to the DM composed of PBHs, in which case there is no need to accumulate DM in order to form a BH. If even a single PBH is captured by a compact star, the latter gets destroyed. Requiring that the probability of such an event is small leads to the constraints on the PBH abundance~\cite{Capela:2013yf,Capela:2012jz}. The key quantity which determines the strength of the constraints is the amount of captured DM. There are two different capture mechanisms. A star can capture DM from its surrounding environment, such as the Galactic halo, during its lifetime. The DM particles passing through the star may interact with the nucleons, losing enough energy to become gravitationally bound~\cite{1995PhRvD..51..328J,2009JCAP...08..037N}. Then each subsequent orbit will also pass through the star, so that eventually, after many collisions, the DM particle will sink to the center of the star. Such capture process can lead to the accumulation of a considerable amount of DM inside the compact star throughout its lifetime~\cite{Kouvaris:2010jy}. The DM could also be captured during the star formation. A main-sequence star is formed from the gravitational collapse of a prestellar core in a giant molecular cloud. In the course of this process, the DM that was initially gravitationally bound to the core undergoes adiabatic contraction, forming a cuspy profile centered at the star, with the density $\rho(r)$ behaving like $\rho(r)\propto r^{-3/2}$. Some of this DM ends up captured inside the star. This mechanism was first discussed in~\cite{Sivertsson:2010zm} and more recently in~\cite{Capela:2012jz} where the constraints on the abundance of PBHs were derived. Note that the adiabatic contraction is a purely gravitational phenomenon that assumes nothing about the DM-to-nucleon cross section~\cite{Gnedin:2004cx}. In this paper we study in more detail the adiabatic contraction of the DM caused by the star formation process, with the final goal to obtain a more precise estimate of the amount of the DM captured by a star as a result of its formation. It turns out that for this purpose the (usually considered) DM density profile after the contraction is not sufficient, and one needs to know in more detail the distribution of DM in the phase space. As we will show, the detailed calculations lead to the results that are qualitatively different from the estimates based on the contracted DM density profile alone. To see what is the point, recall that in order to get captured, the DM particles have to lose their energy by interactions with the star material. Therefore, one needs to calculate the number of particles whose orbits cross the star after the adiabatic contraction. In Ref.~\cite{Capela:2012jz}, this number was estimated by taking the DM density profile after the contraction and integrating it over the volume of the newly-formed star. However, this estimate does not account correctly for particles that spend only a small fraction of time inside the star, because their contribution into the density is correspondingly suppressed. At the same time these particles can still get captured because their orbits will cross the star again and again, and eventually they will get captured if given enough time. The question is how big is the number of such particles. To quantify this effect, we have performed a simulation of the adiabatic contraction process where we have measured, in addition to the density at a given distance, the number of particles that have orbits with the periastron smaller than given radius $r$. Our key observation is that the number of such particles scales differently at small $r$ than the number of particles that {\em are} within $r$ at a given moment: while the latter behaves like $\propto r^{1.5}$ (which corresponds to the density having a spike $\rho\propto r^{-1.5}$), the former goes as $r$. This means that at small $r$ there are much more particles that {\em ever pass} within $r$, than actually {\em are} within $r$ at a given moment. If these particles have enough time to lose their energy and get captured, this would substantially boost the amount of captured DM and improve the constraints accordingly. For relevant values of parameters, the boost factor can be as large as $\sim 2\times 10^3$. The estimate of Ref.~\cite{Capela:2012jz} based on the density profile is, therefore, by far over-conservative. When these considerations are applied to the case of PBH, the previously derived constraints~\cite{Capela:2012jz} get improved. The improvement factor depends on the PBH mass that determines the energy loss time and for some masses can reach its maximum value of $\sim 2\times 10^3$. The rest of this paper is organized as follows. In Sect. II we discuss the capture of DM by adiabatic contraction during the star formation stage. In Sect. III we present the resulting constraints coming from the enhancement of the DM capture in the case the latter is composed of PBHs. Finally, in Sect. IV we present our conclusions. | \label{sec:conclusions} In this paper we have considered the adiabatic contraction of DM during the formation of a star. By simulating the behavior of $\sim 30$ million particles, we reconstructed the phase space distribution of the DM at the end of the star formation. In particular, the number of particles $n(r)$ within a given radius $r$ was found to be proportional to $r^{1.5}$, which corresponds to the DM density profile $\rho(r) \propto r^{-1.5}$, in agreement with previous calculations and the Liouville theorem. At the same time, we have found that the adiabatic contraction creates a rather special distribution of particle orbits. Namely, if one considers the particles that contribute to $n(r)$ for a small $r$, a substantial ($O(1)$) fraction of them have very elongated orbits with periastra smaller than $r$. In fact, the number of particles $\nu(r)$ that have periastra smaller than $r$ scales as $\nu(r) \propto r$. Such particles spend only a small fraction of time close to the center, so their individual contributions to the density at small $r$ are suppressed. However, they are numerous enough to contribute non-negligibly to the density. At $r=R_\odot$, there are about $1.8\times 10^3$ more particles that have periastra smaller than $r$ than there are particles within $r$. This has implications for the DM capture by stars after their formation. A large number of particles that constitute the DM cusp around the newly-formed star have orbits that cross the star, which potentially leads to their capture. This factor has not been taken into account in the previous estimates. As an application, we have considered the capture of PBH by stars, which leads to the constraints on the PBH abundance. We have recalculated the constraints of Ref.~\cite{Capela:2012jz} computing capture times individually for each trajectory, and taking trajectories directly from the simulation of the adiabatic contraction. In the range of PBH masses $10^{20}~\text{g}\lesssim m_{\text{BH}} \lesssim 10^{26}~\text{g}$, the constraints are improved by about 3 orders of magnitude due to the enhancement factor discussed above. The resulting constraints are shown in Fig.~\ref{fig:constraints_compare}. The most stringent constraints are obtained from observations of compact stars in the regions with high DM density and small velocity dispersion. The examples corresponding to the densities $\rho = (10^{4},10^3,10^2)~\text{GeV}~\text{cm}^{-3}$ and a low velocity dispersion $\bar{v} = 7~\text{km}~\text{s}^{-1}$ are shown in Fig.~\ref{fig:constraints_total}. Such conditions could have been present at the cores of metal-poor globular clusters at the epoch of star formation if they are of a primordial origin~\cite{Bertone:2007ae}, or --- with a DM density somewhat smaller than $10^3~\text{GeV}/\text{cm}^3$ --- in dwarf spheroidal galaxies~\cite{Strigari2008,2007PhRvD..75h3526S,2012MNRAS.420.2034S}. Even though clear observations of compact objects in dwarf spheroidals are lacking at the moment, first glimpses of NSs existence may have already been observed~\cite{2013IAUS..291..111R,Maccarone2005}. | 14 | 3 | 1403.7098 |
1403 | 1403.5054_arXiv.txt | We present a new grid of stellar models and nucleosynthetic yields for super-AGB stars with metallicities Z=0.001 and 0.0001, applicable for use within galactic chemical evolution models. Contrary to more metal rich stars where hot bottom burning is the main driver of the surface composition, in these lower metallicity models the effect of third dredge-up and corrosive second dredge-up also have a strong impact on the yields. These metal-poor and very metal-poor super-AGB stars create large amounts of \chem{4}He, \chem{13}C and \chem{14}N, as well as the heavy magnesium isotopes \chem{25}Mg and \chem{26}Mg. There is a transition in yield trends at metallicity Z$\approx$0.001, below which we find positive yields of \chem{12}C, \chem{16}O, \chem{15}N, \chem{27}Al and \chem{28}Si, which is not the case for higher metallicities. We explore the large uncertainties derived from wind prescriptions in super-AGB stars, finding $\approx$ 2 orders of magnitude difference in yields of \chem{22}Ne, \chem{23}Na, \chem{24,25,26}Mg, \chem{27}Al and our s-process proxy isotope $g$. We find inclusion of variable composition low temperature molecular opacities is only critical for super-AGB stars of metallicities below Z$\approx$0.001. We analyze our results, and those in the literature, to address the question: Are super-AGB stars the polluters responsible for extreme population in the globular cluster NGC 2808? Our results, as well as those from previous studies, seem unable to satisfactorily match the extreme population in this globular cluster. | Super-AGB stars are characterised by off-centre carbon ignition and at the low metallicities studied in this work have initial masses between $\sim$ 6.5 and 9.0\,M$_\odot$. They undergo from tens to thousands of thermal pulses and associated third dredge-up (3DU) episodes which enrich their envelopes with the products of nuclear burning. They also have relatively extreme nucleosynthetic conditions with temperatures at the base of the convective envelope reaching over 130MK which leads to efficient hot bottom burning (HBB). Because of their short lifetimes (30-50 Myrs), metal-poor and very metal-poor (Z=0.001 and 0.0001) super-AGB stars are some of the first AGB stars to have enriched the interstellar medium. While nucleosynthesis in intermediate-mass and massive\footnote{We define massive AGB stars as those with initial masses $\ga$ 5\,M$_\odot$ but not massive enough to ignite carbon.} metal-poor and very metal-poor AGB stars has been explored in considerable detail \cite[e.g.][]{ven02,den03,her04a,fen04,ven05,mar09,ven09b,kar10,lug12a,dor13a}, to date, at these metallicities there have only been a few nucleosynthesis studies along the entire super-AGB phase \citep{sie10,ven10a,ven13}, with a notable absence of third dredge-up in these works. At moderate metallicities ([Fe/H] $\ga$ $- 0.6$), massive AGB star models can be directly compared to observations to constrain the occurrence and relative impact of HBB and 3DU on the nucleosynthesis \cite[e.g.][]{ven00,mcs07,van12,gar13}. At the lower metallicities considered here however, super-AGB and massive AGB stars will have died long ago, making direct comparison impossible. Nevertheless there is the possibility to derive some constraints on the evolution and nucleosynthesis of low-metallicity super-AGB and massive AGB star evolution indirectly via galactic chemical evolution models. For example super-AGB and massive AGB stars may have contributed to the rise of heavy magnesium isotopes in the Galaxy \citep{fen03} and may be an extra source of \chem{13}C and \chem{14}N in the early stages of the Galaxy's formation. Super-AGB stars may also play a role as polluters within globular clusters. The ``abundance anomaly problem'' in globular clusters (GCs) has been one of the most thoroughly researched subjects in modern astrophysics \cite[e.g.][]{kra79,kra94,gra12}. A substantial fraction of stars within GCs are found to have unusual compositions (not seen in field stars), characterised by the results of hot hydrogen burning \citep{den90}. These stars show variations in a number of light elements such as C, N, O, F Na, Mg, Al and Si. However, the anomalous and normal stars typically have the same [Fe/H], total number abundance of C+N+O \cite[e.g.][]{smi96,iva99}\footnote{The constancy of total CNO abundance between anomalous and normal stars within GCs is debated, for example in NGC 1851 where it is found to vary by a factor of 2 \citep{yon09} or not at all \citep{vil10}.}, lithium \citep{dor10b} and s-process element abundances \cite[e.g.][]{yon06b,yon08a}. These shared traits place strong constraints on the source of the anomalous material. One of the main theories to explain these abundance anomalies is that GCs are made of multiple generations of stars, with the anomalous stars being formed from the enriched material of a first generation of stars (for a review see \citealt{gra12}). Stars within GCs are also further divided into three main populations based on their O and Na abundances; primordial (P) which comprises of first generation stars and intermediate (I) and extreme (E) which belong to the second generation. Generally only the most massive CGs harbour an extreme population \citep{car09a} showing the largest abundance anomalies, with substantial depletion of C, O, Mg and large enhancement of He, N, Na and Al. Although the source of gas from which the suspected second generation of stars formed is still uncertain, proposed candidates include: intermediate-mass (and super-)AGB stars \cite[]{cot81,ven01,der08}, winds from fast rotating massive stars \cite[]{nor04,dec07b}, massive star binaries \citep{dem09} and super-massive stars \citep{den14} to name a few. While each of these classes of polluter has their associated problems matching certain observed features \citep{gra12} here we focus on the super-AGB and massive AGB star scenario. \cite{der08,der12} have suggested the extreme population within GCs are formed directly from pristine super-AGB ejecta, with the intermediate population formed via some dilution of polluted material with pristine gas \cite[e.g.][]{bek07,dan07,der08,der10,gra10}. This paper is the third in this series dedicated to the study of super-AGB stars and is organized as follows: Section~\ref{sec-code} summarizes our numerical program and input physics, in Section~\ref{sec-nucleo} we explore the nucleosynthesis of these low metallicity models, in Section~\ref{yields} we discuss the stellar yields, and a range of uncertainties. Comparisons to the past works on super-AGB nucleosynthesis are made in Section~\ref{sec-ngc2808} and then we apply these yield results to the extreme population in the globular cluster NGC 2808. In Section~\ref{sec-conclude} we summarize and conclude. | We have computed a grid of metal-poor and very metal-poor super-AGB stars to explore element production. These stars create large amounts of \chem{4}He, \chem{13}C, \chem{14}N and \chem{17}O, as well as the heavy magnesium isotopes \chem{25}Mg and \chem{26}Mg. In addition, and contrary to higher metallicity models, we also find positive yields of \chem{12}C, \chem{15}N, \chem{16}O, and the heavier proton chain species \chem{27}Al and \chem{28}Si. The occurrence of third dredge-up in our models is a key difference compared to previous low metallicity super-AGB yield studies. Whilst there is evidence for third dredge-up in intermediate-mass/massive AGB stars of higher metallicities (e.g. Rb observations - \citealt{gar06,gar09}, very luminous C-stars - \citealt{van99}, large N overabundance - \citealt{mcs07}), at the metallicities considered here, the efficiency, or even the occurrence of 3DU is unknown. Contrary to the higher metallicity models where the surface composition is driven almost purely by HBB, the pollution from dredge-up events plays an important role in our metal-poor and very metal-poor models. The dominant dredge-up process to affect the surface enrichment changes from 3DU at $\sim$ 6.5\,M$_\odot$, to about equal contribution from C2DU and 3DU at $\sim$ 7.0\,M$_\odot$ whilst the abundances in the more massive models are mainly dictated by the large enrichment prior to the TP-(S)AGB phase from CO2DU/dredge-out. The envelope enrichment in carbon from either corrosive 2DU or a dredge-out event prior to the TP-(S)AGB phase can have a large impact on the further evolution, primarily though enhanced mass loss driven via enhanced opacity. We have explored a selection of uncertainties within our super-AGB star models. The mass-loss rate has the greatest impact with up to a factor of 10 difference in TP-(S)AGB lifetime of very metal-poor stars when two commonly used mass-loss rate prescriptions are used. The yields of isotopes most affected by changes to the mass-loss prescriptions are \chem{22}Ne, \chem{23}Na and \chem{24,25,26}Mg and \chem{27}Al and $g$. The appropriate treatment of low temperature molecular opacities is crucial to super-AGB star evolution for metallicities below Z$\approx$0.001, where they truncate the evolution by approximately 50 per cent. However, the inclusion of these updated opacities has a less dramatic effect than changes to the mass-loss rate prescription. A modest increase in $\alpha_{\rm{mlt}}$ teamed with a rapid mass-loss rate was found to have only a small impact on super-AGB star nucleosynthesis. Extrapolated thermal pulses at the end of the evolution, to account for possible nucleosynthesis after convergence issues cease calculations, have negligible effect on the yields of most isotopes. With the yields showing \textit{such} large variations both within, and between, different research groups, which observable nucleosynthetic signature of low metallicity super-AGB stars can be considered robust? As the CNO elemental yields are dependent on 3DU efficiency and mass-loss prescription they show no consistency between code results. Lithium is notoriously temperamental, and is highly dependent on the mass-loss rate and the treatment of convective mixing. The Ne-Na cycle and Mg-Al chain isotope yields vary widely due to uncertainties in reaction rates, mass-loss prescriptions and occurrence of 3DU. Also with large variation in 3DU efficiency between calculations (0 $\la$ $\lambda$ $\la$ 1), all elements heavier than iron will vary considerable between calculations. This leaves the large helium enhancement from 2DU, as well as certain isotopic ratios, such as \chem{12}C/\chem{13}C and \chem{14}N/\chem{15}N which reach their equilibrium values as the only consistent result between different studies. However, these model independent characteristics are not unique to super-AGB stars but are also found in intermediate-mass and massive AGB stars. Whilst super-AGB are thought of as one of the most likely candidate polluters of the extreme population within NGC 2808, Fig.~\ref{fig-giant} shows that the observational data for this cluster and theoretical super-AGB yield predictions are not compatible. For the super-AGB pollution scenario to be salvaged extra-mixing within the extreme population needs to be invoked, and some observational results would have to be questioned. However, given the large quantitative uncertainties in both the theoretical and observational results we cannot completely rule out super-AGB stars. | 14 | 3 | 1403.5054 |
1403 | 1403.2218_arXiv.txt | {The Cherenkov Telescope Array (CTA) will be the world's first observatory for detecting gamma-rays from astrophysical phenomena and is now in its prototyping phase with construction expected to begin in 2015/16. In this work we present the results from early attempts at detailed simulation studies performed to assess the need for atmospheric monitoring. This will include discussion of some lidar analysis methods with a view to determining a range resolved atmospheric transmission profile. We find that under increased aerosol density levels, simulated gamma-ray astronomy data is systematically shifted leading to softer spectra. With lidar data we show that it is possible to fit atmospheric transmission models needed for generating lookup tables, which are used to infer the energy of a gamma-ray event, thus making it possible to correct affected data that would otherwise be considered unusable.} | \noindent The Cherenkov Telescope Array (CTA) aims to increase its flux sensitivity by an order of magnitude compared to existing ground-based gamma-ray telescopes \cite{bib:mcpaper}. This will be achieved using Cherenkov telescopes of 3 different sizes, a large size telescope (LST) $\mathrm{\sim 23\thickspace m}$ diameter, a medium size telescope (MST) $\mathrm{\sim 12\thickspace m}$ diameter and a small size telescope (SST) $\mathrm{\sim 4\thickspace m}$ diameter. In order to achieve such sensitivity gains it is important to minimise the systematic uncertainty in derived flux and energy resolution. Imaging atmospheric Cherenkov telescopes are calorimetric by nature and as such a good knowledge of atmospheric conditions is required at the telescope site. Atmospheric quality affects both shower development and the Cherenkov yield in two ways. Firstly, in the production of Cherenkov light atmospheric quality affects the vertical profile of the refractive index of the air and hence shower development. This variation is seasonal, and effects mid-latitudes worse than the tropics. However, the profile can be measured using radiosondes for example and any seasonal variation that might exist can be accounted for. It is also possible for high-level aerosols (e.g. clouds) to occur around shower maximum and so affect Cherenkov yield and image shape. \myparskip \noindent Secondly, poor atmospheric quality can also result in the loss of Cherenkov light. For example atmospheric quality affects Cherenkov light propagation through Rayleigh \& Mie scattering of the Cherenkov light, which can lower the brightness of an image in the camera for a shower of given energy and core distance. However, by using lidar measurements it is possible to derive a range-resolved probability of transmission (at the lidar wavelength) and adjust atmospheric models, needed in simulations used to reconstruct gamma-ray spectra, accordingly \cite{bib:nolan}. \myparskip \noindent This work highlights an early simulation study conducted using a hypothetical 97 telescope array to illustrate the effects of atmospheric quality on a reconstructed gamma-ray spectrum. \myparskip \noindent Finally, another motivational factor for a large observatory like CTA is the desire to maximise the duty cycle of the instrument. Thus being able to resurrect otherwise unusable data due to relatively poor atmospheric conditions becomes important. | \noindent Currently within the field of ground-based gamma-ray astronomy, atmospheric quality is accounted for by monitoring the background cosmic-ray trigger rates and data with sub-standard atmospheric quality is discarded \cite{bib:nolan}. This work shows that it is possible to use a lidar to take in-situ atmospheric measurements in order to derive the probability of transmission at a wavelength close to the maximum for Cherenkov light production. A model of atmospheric transmission for a spectrum of wavelengths is then fitted to the lidar data and used within simulations to produce lookup tables that better reflect the actual atmospheric quality. Correcting for changing atmospheric quality in such a way can increase the lifetime of an observatory like CTA. In addition, such an active atmospheric calibration method helps to lower any systematic uncertainty on the derived flux. \myparskip \par{\noindent {\it Acknowledgements} The authors would like to acknowledge the support of their host institutions and also the support from the Science and Technology Facilities Council of the UK. In addition, the authors gratefully acknowledge support from the agencies and organizations listed on this page: http://www.cta-observatory.org/?q=node/22} | 14 | 3 | 1403.2218 |
1403 | 1403.5324_arXiv.txt | The measurement of B-mode polarization of the cosmic microwave background at large angular scales by the BICEP experiment suggests a stochastic gravitational wave background from early-universe inflation with a surprisingly large amplitude. The power spectrum of these tensor perturbations can be probed both with further measurements of the microwave background polarization at smaller scales, and also directly via interferometry in space. We show that sufficiently sensitive high-resolution B-mode measurements will ultimately have the ability to test the inflationary consistency relation between the amplitude and spectrum of the tensor perturbations, confirming their inflationary origin. Additionally, a precise B-mode measurement of the tensor spectrum will predict the tensor amplitude on solar system scales to 20\% accuracy for an exact power law tensor spectrum, so a direct detection will then measure the running of the tensor spectral index to high precision. | 14 | 3 | 1403.5324 |
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1403 | 1403.2610.txt | When comet nuclei approach the Sun, the increasing energy flux through the surface layers leads to sublimation of the underlying ices and subsequent outgassing that promotes the observed emission of gas and dust. While the release of gas can be straightforwardly understood by solving the heat-transport equation and taking into account the finite permeability of the ice-free dust layer close to the surface of the comet nucleus, the ejection of dust additionally requires that the forces binding the dust particles to the comet nucleus must be overcome by the forces caused by the sublimation process. This relates to the question of how large the tensile strength of the overlying dust layer is. Homogeneous layers of micrometer-sized dust particles reach tensile strengths of typically $10^3$ to $10^4$ Pa. This exceeds by far the maximum sublimation pressure of water ice in comets. It is therefore unclear how cometary dust activity is driven. \par To solve this paradox, we used the model by Skorov and Blum (Icarus 221, 1-11, 2012), who assumed that cometesimals formed by gravitational instability of a cloud of dust and ice aggregates and calculated for the corresponding structure of comet nuclei tensile strength of the dust-aggregate layers on the order of 1 Pa. Here we present evidence that the emitted cometary dust particles are indeed aggregates with the right properties to fit the model by Skorov and Blum. Then we experimentally measure the tensile strengths of layers of laboratory dust aggregates and confirm the values derived by the model. To explain the comet activity driven by the evaporation of water ice, we derive a minimum size for the dust aggregates of $\sim 1$ mm, in agreement with meteoroid observations and dust-agglomeration models in the solar nebula. Finally we conclude that cometesimals must have formed by gravitational instability, because all alternative formation models lead to higher tensile strengths of the surface layers. | Introduction: formation scenarios of planetesimals and cometesimals} It is now well established that dust inside the snow line of the solar nebula quickly coagulated into millimeter- to centimeter-sized agglomerates due to direct sticking in collisions \citep{Guettler2010,Zsom2010}. The further growth to planetesimal-sized objects is still under debate, with two major scenarios under consideration: the mass transfer scenario (1) and the gravitational instability scenario (2). \par (1) As direct sticking is mostly prevented by bouncing among the dust aggregates \citep{BlumMuench1993,Langkowski2008,Weidling2009,Beitz2012,Weidling2011,Schraepler2012,Deckers2013}, only those particles colliding with velocities slower than the sticking-bouncing transition can further grow, whereas the fastest collisions in the ensemble lead to fragmentation with mass transfer \citep{Windmark2012,Windmark2012b,Garaud2013}. This latter process has been extensively studied in the laboratory \citep{Wurm2005,Teiseretal2009a,Teiseretal2009b,Guettler2010,Kothe2010,Teiseretal2011} and is now well established for aggregates consisting of micrometer-sized silicate grains. It has been shown that in principle planetesimals can form by this process \citep{Windmark2012,Windmark2012b,Garaud2013} although the timescales are rather long and details about counteracting processes \citep[e.g., erosion;][]{Schraepler2011} need to be clarified. \par (2) Alternatively, \citet{Johansen2007} showed that cm-sized particles can be sufficiently concentrated by the streaming instability \citep{Youdin2005} so that the ensemble becomes gravitationally unstable and forms planetesimals. Also here, several details need to be clarified before this process can be regarded as established, e.g. the collisional fate of the dust agglomerates within the instabilities, fragmentation of the collapsing cloud and the mass distribution function of the resulting planetesimals, and the required high metallicity of the solar nebula. \par In the outer solar nebula beyond the snow line, the dominant material should be (water) ice. Due to the higher anticipated stickiness of water-ice particles \citep{Gundlach2011b}, ice aggregates are supposed to grow to larger masses and fluffier structures in the outer solar nebula \citep{Wada2008,Wada2009,Okuzmi2012,Kataokaetal2013}. As empirical proof for this concept from laboratory experiments is still missing, it can only be speculated whether icy planetesimals form directly by hit-and-stick collisions, or by a multi-step process. If direct formation of cometesimals is not feasible, similar processes as discussed above for the inner solar nebula might also apply for its outer reaches. \par Since the first space missions to comet Halley it has been known that comets consist in almost equal parts of ice and refractory materials (dust), with the addition of organic materials \citep{Jessbergeretal1988}, which in turn led to revised cometary dust modeling \citep{Greenberg1990,Li1997,Greenberg1998}. The samples brought back from comet 81P/Wild by the Stardust mission revealed that the refractory materials are high-temperature condensates, which must have been radially mixed outwards before the formation of the comet nucleus \citep{McKeegan2006,Zolensky2006}. As the growth timescales to mm or cm sizes are rather short in the inner solar nebula (a few $10^3$ years) and as the dust aggregates are supposed to be rather compact (with a porosity of ``only'' $60-65 \, \%$, according to \citet{Zsom2010} and \citet{Weidling2009}), with any further growth slowed down due to the decreased stickiness of large dust aggregates \citep{Guettler2010}, it is plausible to assume that the refractory materials were mixed into the outer solar nebula in form of mm- to cm-sized agglomerates (see also Sects. \ref{sec:comcomp} and \ref{sec:comtensile}). Hence, cometesimals in the outer solar nebula were then formed out of icy and dusty agglomerates by one of the two processes described above, namely (1) fragmentation with mass transfer (MT) or (2) spatial enhancements in (magneto-)hydrodynamic instabilities with subsequent gravitational instability (GI). From this line of reasoning, we can derive several physical distinctions in the resulting icy-dusty planetesimals. We summarize these in the Table \ref{tab:comprop}. \begin{table*} \centering \caption{Comparison between the two formation scenarios of icy-dusty planetesimals. GI stands for gravitational instability, MT represents the process of mass transfer.}\label{tab:comprop} \scriptsize %\small \vspace{2mm} \begin{tabular}{lcccc} \hline & GI & & MT & \\ \hline Volume filling factor & $0.35 \times 0.6 \approx 0.2$ &[1,7] & $\sim$0.4 &[2] \\ Tensile strength of interior [Pa] & $\sim 10$ &[3] & $\sim 10000$ &[2,5] \\ Tensile strength of ice-free outer dust layer [Pa] & $\sim 1$ &[3] & $\sim 1000$ &[2,5] \\ Gas permeability [$\rm m^4 s^{-1}$] & $\sim 1 \times 10^{-6}$ &[4] & $\sim 1 \times 10^{-9}$& [4] \\ Thermal conductivity [$\rm W m^{-1} K^{-1}$] & $10^{-3}-1$ &[6] & $10^{-2}-10^{-1}$& [6] \\ & (conduction/radiation)& & (conduction) &\\ \hline \multicolumn{4}{l}{References:}\\[-0.2cm] \multicolumn{4}{l}{[1] \citet{Weidling2009}, [2] \citet{Kothe2010}, [3] \citet{Skorov2012}, [4] \citet{Gundlach2011}, [5] \citet{Blum2006a},}\\[-0.2cm] \multicolumn{4}{l}{[6] \citet{Gundlach2012}, [7] \citet{Zsom2010}.} \end{tabular} \end{table*} It should be mentioned that we assume that the formation process for cometesimals was the same anywhere in the outer solar nebula. Thus, the following discussion in this paper refers to both, Kuiper-belt and Oort-cloud comets. As to the formation timescales for cometesimals, these are required to be long enough for the radial mixing of the high-temperature condensates to occur, but certainly shorter that the lifetime of the nebula gas. As this might be a problem for the MT origin of cometesimals at large heliocentric distances, the timescales for the instability-driven formation of cometesimals should always be sufficiently short. In the latter, however, the aggregate sizes at which the bouncing barrier is reached and for which then some concentration process forms a gravitationally unstable cloud, could be considerably different (albeit yet unknown) for the two reservoirs of Kuiper-belt and Oort-cloud comets. \par The volume filling factor $\phi$ is defined as the fraction of the total volume occupied by the material and is related to the porosity $\psi$ by $\phi \, = \, 1 \, -\, \psi$. For an icy-dusty planetesimal formed by the GI process, the volume filling factor is determined by the packing fraction of the dust aggregates into the planetesimal ($\phi_{\rm global} \approx 0.6$, if we assume that the dust aggregates pack almost as densely as possible) and by the volume filling factor of the individual dust aggregates \citep[$\phi_{\rm local} \approx 0.35$, according to][]{Weidling2009}. The volume filling factor of the MT dust aggregates has been measured to be close to $\phi_{\rm local} = 0.4$ \citep{Kothe2010}. The tensile strength of a package of dust aggregates, which collapsed under their own gravity to form a km-sized body with a volume filling factor of $\phi_{\rm global}$ has been calculated by \citet{Skorov2012} to be \begin{equation}\label{Eq:tsmodel} p_{\rm tensile} \, = \, p'_{\rm tensile} \phi_{\rm global} \left( \frac{s}{\mathrm{1 mm}} \right)^{-2/3} , \end{equation} with $p'_{\rm tensile}=1.6 \, \mathrm{Pa}$ and $s$ being the radius of the infalling dust aggregates. For ice aggregates, the tensile strength is supposed to be a factor of ten higher \citep{Gundlach2011b}. In the case of planetesimals formed by the MT process, their rather compact packing of the monomer grains ensures a relatively higher tensile strength of $\sim 1 \, \mathrm{kPa}$ for volatile-free and $\sim 10 \, \mathrm{kPa}$ for icy particles \citep{Blum2006a}. Due to the smaller pore size in the planetesimals formed by MT (the pore size is on the order of the monomer-grain size, i.e., $\sim 1 \, \mathrm{\mu m}$, whereas for planetesimals formed by GI the pore size is on the order of the aggregate size), the gas permeability is much lower \citep{Gundlach2011}. The thermal conductivity is not easily distinguishable between the two formation models, due to the fact that for large pore sizes, radiative energy transport is no longer negligible \citep{Gundlach2012}. Thus, the range of possible thermal-conductivity values is much larger for the GI-formed planetesimals than for those formed by MT. \par As mentioned above, \citet{Skorov2012} were the first to bring up the distinction in tensile strength between the two models, who related the formation of icy-dusty planetesimals to present comet nuclei, and who showed that, according to their model (see Eq. \ref{Eq:tsmodel}), only the GI model can explain a continued gas and dust activity of a comet. Their model for the tensile strength of the ice-free outer layers of a comet nucleus is based on the assumption that dust and ice aggregates once formed the comet nucleus by gravitational instability so that essentially the aggregates collapsed below or at the very low escape speed of the kilometer-sized body. Thus, the aggregates are only slightly deformed and the resulting binding between the clumps is much weaker than in the mass-transfer process. \par In this article, we intend to verify the model by \citet{Skorov2012} and to support their statement that comets were formed in gravitational instabilities. This will be done in the following: in Sect. \ref{sec:comcomp} we will show that comet nuclei indeed consist of mm- to cm-sized dust particles and ice clumps with at least these sizes. In Sect. \ref{sec:comtensile}, we will then show that observed cometary dust aggregates or meteoroids are consistent with model expectations for dust aggregates in the bouncing regime, i.e. with a rather low porosity and a correspondingly rather low tensile strength $p_{\rm tensile}$ (in this paper, we denote $p_{\rm tensile} \sim 1~\mathrm{kPa}$ as low tensile strength and $p_{\rm tensile} \sim 1~\mathrm{Pa}$ as ultra-low tensile strength)\footnote{Here, it should be mentioned that the internal tensile strength of aggregates and meteoroids is low ($p_{\rm tensile} \sim 1-10~\mathrm{kPa}$), whereas an arrangement of aggregates possess an ultra-low tensile strength ($p_{\rm tensile} \sim 1~\mathrm{Pa}$).}. In the following experimental part (see Sect. \ref{sec:lab}), we will then construct an analog sample of a cometary ice-free surface and measure its tensile strength. The results of our laboratory experiments and their analysis will then be applied to comet nuclei (see Sect. \ref{sec:application}). We conclude in Sect. \ref{sec:conclusion} that comets were formed by gravitational instability and will give a short outlook of future work in Sect. \ref{sec:future}. | Conclusion} A wealth of laboratory and microgravity experiments on dust-aggregate collisions have led to a collision model for protoplanetary dust \citep{Guettler2010}. With this dust-aggregate collision model, the growth of millimeter to centimeter-sized dust aggregates in protoplanetary disks can be explained. Ice should behave qualitatively similar, but quantitative differences are expected because (i) ice is stickier than silicates \citep{Gundlach2011b} and (ii) the collision velocities are different at 30 AU \citep{Weidenschilling1997}. As we do not have an analog of the dust-aggregate collision model for ice aggregates, it is unclear whether the expected sizes of the ice aggregates exceed those of the dust aggregates or not. For our model, this is, however, of minor importance, because we assume that dust and ice aggregates are separated inside the comet and the ice sublimation leaves the dust aggregates intact. This is consistent with meteoroid survival in interplanetary space for thousands of years \citep{Jenniskens1998,Sykesetal2004,Trigoetal2005}. However, in our model the ice-aggregate sizes cannot exceed the dust-aggregate sizes by much, because after ice sublimation, a dust layer of thickness of the ice-aggregates forms, which, once closed, might be too thick to sustain cometary activity if the ice aggregates were too big (see Sect. \ref{sec:application}). \par As was shown in Sect. \ref{sec:introduction}, models of the formation of planetesimals/cometesimals require either collective particle effects (e.g., the streaming instability and gravitational instability), a continuous fragmentation-agglomeration cycle, or more sticky materials than silicates (e.g., ices or aggregates-of-aggregates). If comet nuclei are leftover cometesimals, whose formation is described by one of the two competing planetesimal-formation models (gravitational instability or mass transfer), then only very gentle processes, i.e. the gravitational-instability model, can explain the gas and dust activity (i.e. the very small required tensile strengths) of comets. The mass-transfer model predicts comets with much too high tensile strengths. For the gravitational-instability formation scenario of cometesimals, the model by \citep{Skorov2012} predicts the correct tensile strength as proven by our experiments (see Sect. \ref{sec:lab}). Our experiments confirm that if the comet nuclei consist of dust (and ice) aggregates, the cohesion-induced tensile strength of the ice-free surface layer is $p_{\rm tensile} \sim 1 ~ \mathrm{Pa}~ (s/{\mathrm{1 mm}})^{-\alpha}$, with $s$ being the radius of the dust aggregates and $\alpha = 2/3$. \par Our experiments also show that hydrostatic compression of layers of dust aggregates leads to an additional consolidation and, thus, an increase of tensile strength towards the comet-nucleus center (see Sect. \ref{sec:lab}). Impact-induced consolidation \citep{Beitzetal2013} can also lead to an increased tensile strength. Both processes might explain the extinction of old comets \citep{Jewitt2008}, if there is no re-organization of dust aggregates after the ice has evaporated. | 14 | 3 | 1403.2610 |
1403 | 1403.0007.txt | Over the past two decades, an avalanche of data from multiwavelength imaging and spectroscopic surveys has revolutionized our view of galaxy formation and evolution. Here we review the range of complementary techniques and theoretical tools that allow astronomers to map the cosmic history of star formation, heavy element production, and reionization of the Universe from the cosmic ``dark ages'' to the present epoch. A consistent picture is emerging, whereby the star-formation rate density peaked approximately 3.5 Gyr after the Big Bang, at $z \approx 1.9$, and declined exponentially at later times, with an e-folding timescale of 3.9 Gyr. Half of the stellar mass observed today was formed before a redshift $z = 1.3$. About 25\% formed before the peak of the cosmic star-formation rate density, and another 25\% formed after $z = 0.7$. Less than $\sim 1$\% of today's stars formed during the epoch of reionization. Under the assumption of a universal initial mass function, the global stellar mass density inferred at any epoch matches reasonably well the time integral of all the preceding star-formation activity. The comoving rates of star formation and central black hole accretion follow a similar rise and fall, offering evidence for co-evolution of black holes and their host galaxies. The rise of the mean metallicity of the Universe to about 0.001 solar by $z = 6$, one Gyr after the Big Bang, appears to have been accompanied by the production of fewer than ten hydrogen Lyman-continuum photons per baryon, a rather tight budget for cosmological reionization. | \label{sec:introduction} The origin and evolution of galaxies are among the most intriguing and complex chapters in the formation of cosmic structure, and observations in this field have accumulated at an astonishing pace. Multiwavelength imaging surveys with the {\it Hubble} (HST) and {\it Spitzer} space telescopes and ground-based facilities, together with spectroscopic follow-up with 8-m-class telescopes, have led to the discovery of galaxies with confirmed redshifts as large as $z = 7.5$ (Finkelstein \etal 2013), as well as compelling photometric candidates as far back as $z \approx 11$ (Coe \etal 2013) when the Universe was only 3\% of its current age. Following the seminal work of Steidel \etal (1995), color-selection criteria that are sensitive to the presence of intergalactic \HI\ absorption features in the spectral energy distribution (SED) of distant sources have been used to build increasingly large samples of star-forming galaxies at $2.5\lta z \lta 9$ (e.g., Madau \etal 1996, Steidel \etal 2003, Giavalisco \etal 2004a, Bouwens \etal 2011b). Infrared (IR)-optical color selection criteria efficiently isolate both actively star-forming and passively evolving galaxies at $z \approx 2$ (Franx \etal 2003, Daddi \etal 2004). Photometric redshifts have become an unavoidable tool for placing faint galaxies onto a cosmic timeline. {\it Spitzer,} {\it Herschel}, and submillimeter telescopes have revealed that dusty galaxies with star-formation rates (SFRs) of order 100$\,\msun$~year$^{-1}$ or more were abundant when the Universe was only 2--3 Gyr old (Barger \etal 1998, Daddi \etal 2005, Gruppioni \etal 2013). Deep near-infrared (NIR) observations are now commonly used to select galaxies on the basis of their optical rest-frame light and to chart the evolution of the global stellar mass density (SMD) at $0<z<3$ (Dickinson \etal 2003). The {\it Galaxy Evolution Explorer} (GALEX) satellite has quantified the ultraviolet galaxy luminosity function (LF) of galaxies in the local Universe and its evolution at $z\lta 1$. Ground-based observations and, subsequently, UV and IR data from {GALEX} and {Spitzer} have confirmed that star-formation activity was significantly higher in the past (Lilly \etal 1996, Schiminovich \etal 2005, Le~Floc'h \etal 2005). In the local Universe, various galaxy properties (colors, surface mass densities, and concentrations) have been observed by the {Sloan Digital Sky Survey} (SDSS) to be ``bimodal'' around a transitional stellar mass of $3\times 10^{10}\,\msun$ (Kauffmann \etal 2003), showing a clear division between faint, blue, active galaxies and bright, red, passive systems. The number and total stellar mass of blue galaxies appear to have remained nearly constant since $z\sim 1$, whereas those of red galaxies (around $L^\ast$) have been rising (Faber \etal 2007). At redshifts $0 < z < 2$ at least, and perhaps earlier, most star-forming galaxies are observed to obey a relatively tight ``main-sequence'' correlation between their SFRs and stellar masses (Brinchmann \etal 2004, Noeske \etal 2007, Elbaz \etal 2007, Daddi \etal 2007). A minority of starburst galaxies have elevated SFRs above this main sequence as well as a growing population of quiescent galaxies that fall below it. With the avalanche of new data, galaxy taxonomy has been enriched by the addition of new acronyms such as LBGs, LAEs, EROs, BzKs, DRGs, DOGs, LIRGs, ULIRGs, and SMGs. Making sense of it all and fitting it together into a coherent picture remains one of astronomy's great challenges, in part because of the observational difficulty of tracking continuously transforming galaxy sub-populations across cosmic time and in part because theory provides only a partial interpretative framework. The key idea of standard cosmological scenarios is that primordial density fluctuations grow by gravitational instability driven by cold, collisionless dark matter, leading to a ``bottom-up'' $\Lambda$CDM (cold dark matter) scenario of structure formation (Peebles 1982). Galaxies form hierarchically: Low-mass objects (``halos'') collapse earlier and merge to form increasingly larger systems over time -- from ultra-faint dwarfs to clusters of galaxies (Blumenthal \etal 1984). Ordinary matter in the Universe follows the dynamics dictated by the dark matter until radiative, hydrodynamic, and star-formation processes take over (White \& Rees 1978). The ``dark side'' of galaxy formation can be modeled with high accuracy and has been explored in detail through $N$-body numerical simulations of increasing resolution and size (e.g., Davis \etal 1985, Dubinski \& Carlberg 1991, Moore \etal 1999, Springel \etal 2005, 2008, Diemand \etal 2008, Stadel \etal 2009, Klypin \etal 2011). However, the same does not hold for the baryons. Several complex processes are still poorly understood, for example, baryonic dissipation inside evolving CDM halos, the transformation of cold gas into stars, the formation of disks and spheroids, the chemical enrichment of gaseous material on galactic and intergalactic scales, and the role played by ``feedback'' [the effect of the energy input from stars, supernovae (SNe), and massive black holes on their environment] in regulating star formation and generating galactic outflows. The purely phenomenological treatment of complex physical processes that is at the core of semi-analytic schemes of galaxy formation (e.g., White \& Frenk 1991, Kauffmann \etal 1993, Somerville \& Primack 1999, Cole \etal 2000) and -- at a much higher level of realism -- the ``subgrid modeling'' of star formation and stellar feedback that must be implemented even in the more accurate cosmological hydrodynamic simulations (e.g., Katz \etal 1996, Yepes \etal 1997, Navarro \& Steinmetz 2000, Springel \& Hernquist 2003, Keres al. 2005, Ocvirk \etal 2008, Governato \etal 2010, Guedes \etal 2011, Hopkins \etal 2012, Kuhlen \etal 2012, Zemp \etal 2012, Agertz \etal 2013) are sensitive to poorly determined parameters and suffer from various degeneracies, a weakness that has traditionally prevented robust predictions to be made in advance of specific observations. Ideally, an in-depth understanding of galaxy evolution would encompass the full sequence of events that led from the formation of the first stars after the end of the cosmic dark ages to the present-day diversity of forms, sizes, masses, colors, luminosities, metallicities, and clustering properties of galaxies. This is a daunting task, and it is perhaps not surprising that an alternative way to look at and interpret the bewildering variety of galaxy data has become very popular in the past two decades. The method focuses on the emission properties of the galaxy population as a whole, traces the evolution with cosmic time of the galaxy luminosity density from the far-UV (FUV) to the far-infrared (FIR), and offers the prospect of an empirical determination of the global history of star formation and heavy element production of the Universe, independently of the complex evolutionary phases of individual galaxy subpopulations. The modern version of this technique relies on some basic properties of stellar populations and dusty starburst galaxies: \begin{enumerate} \item The UV-continuum emission in all but the oldest galaxies is dominated by short-lived massive stars. Therefore, for a given stellar initial mass function (IMF) and dust content, it is a direct measure of the instantaneous star-formation rate density (SFRD). \item The rest-frame NIR light is dominated by near-solar-mass evolved stars that make up the bulk of a galaxy's stellar mass and can then be used as a tracer of the total SMD. \item Interstellar dust preferentially absorbs UV light and re-radiates it in the thermal IR, so that the FIR emission of dusty starburst galaxies can be a sensitive tracer of young stellar populations and the SFRD. \end{enumerate} By modeling the emission history of all stars in the Universe at UV, optical, and IR wavelengths from the present epoch to $z\approx 8$ and beyond, one can then shed light on some key questions in galaxy formation and evolution studies: Is there a characteristic cosmic epoch of the formation of stars and heavy elements in galaxies? What fraction of the luminous baryons observed today were already locked into galaxies at early times? Are the data consistent with a universal IMF? Do galaxies reionize the Universe at a redshift greater than 6? Can we account for all the metals produced by the global star-formation activity from the Big Bang to the present? How does the cosmic history of star formation compare with the history of mass accretion onto massive black holes as traced by luminous quasars? This review focuses on the range of observations, methods, and theoretical tools that are allowing astronomers to map the rate of transformation of gas into stars in the Universe, from the cosmic dark ages to the present epoch. Given the limited space available, it is impossible to provide a thorough survey of such a huge community effort without leaving out significant contributions or whole subfields. We have therefore tried to refer only briefly to earlier findings, and present recent observations in more detail, limiting the number of studies cited and highlighting key research areas. In doing so, we hope to provide a manageable overview of how the field has developed and matured in line with new technological advances and theoretical insights, and of the questions with which astronomers still struggle nowadays. The remainder of this review is organized as follows. The equations of cosmic chemical evolution that govern the consumption of gas into stars and the formation and dispersal of heavy elements in the Universe as a whole are given in Section \ref{sec:chemev}. We turn to the topic of measuring mass from light, and draw attention to areas of uncertainty in Section \ref{sec:massfromlight}. Large surveys, key data sets and the analyses thereof are highlighted in Section \ref{sec:surveys}. An up-to-date determination of the star-formation history (SFH) of the Universe is provided and its main implications are discussed in Serction \ref{sec:obs_to_para}. Finally, we summarize our conclusions in Section \ref{sec:conclusion}. Unless otherwise stated, all results presented here will assume a ``cosmic concordance cosmology'' with parameters $(\Omega_M, \Omega_\Lambda, \Omega_b, h)=(0.3, 0.7,$ $0.045, 0.7)$. | \label{sec:conclusion} The cosmic history of star formation is one of the most fundamental observables in astrophysical cosmology. We have reviewed the range of complementary techniques and theoretical tools that are allowing astronomers to map the transformation of gas into stars, the production of heavy elements, and the reionization of the Universe from the cosmic dark ages to the present epoch. Under the simple assumption of a universal IMF, there is reasonable agreement between the global SMD inferred at any particular time and the time integral of all the preceding instantaneous star-formation activity, although modest offsets may still point toward systematic uncertainties. A consistent picture is emerging, whereby the SFRD peaked $\sim $3.5 Gyr after the Big Bang, and dropped exponentially at $z<1$ with an e-folding timescale of 3.9 Gyr. The Universe was a much more active place in the past: Stars formed at a peak rate approximately nine times higher than is seen today. Approximately 25\% of the present-day SMD formed at $z > 2$, before the peak of the SFRD, and another 25\% formed since $z = 0.7$, i.e., roughly over the last half of the Universe's age. From the peak of the SFRD at $z \approx 2$ to the present day, and perhaps earlier as well, most stars formed in galaxies that obey a relatively tight SFR--$M_\ast$ correlation, and only a small fraction formed in starbursts with significantly elevated specific SFRs. The smooth evolution of this dominant main-sequence galaxy population suggests that the evolution of the cosmic SFH is primarily determined by a balance between gas accretion and feedback processes, both closely related to galaxy mass, and that stochastic events such as merger-driven starbursts play a relatively minor role. The growth histories of the stellar component of galaxies and their central black holes are similar in shape, suggesting broad co-evolution of black holes and their host galaxies. The rise of the mean metallicity of the Universe to 0.001 solar by redshift six, 1 Gyr after the Big Bang, appears to have been accompanied by the production of fewer than 10 hydrogen LyC photons per baryon, indicating a rather tight budget for cosmological reionization. The SFRD at $z \approx 7$ was approximately the same as that of today, at $z \approx 0$, but only 1\% of today's SMD was formed during the epoch of reionization. As far as the observations and data are concerned, there is still room for improvement in both SFRD and SMD measurements at virtually every redshift, from the local Universe to the epoch of reionization (Section \ref{sec:stateoftheart}). That said, it would be somewhat surprising if new {measurements} changed the picture dramatically at $z < 1$; it is more likely that stellar population modeling, e.g., for deriving stellar masses or SFRs, could still change the details of the picture during the decline and fall of the cosmic SFH. Indeed, at all redshifts, limitations of our methods for interpreting light as mass may play a significant, even dominant, role in the error budget for the analyses described in this review. The peak era of cosmic star formation has been extensively mapped, and yet even with the current data (Figure ~\ref{fig9}), it is still hard to accurately pinpoint the redshift of maximum SFRD within a range $\Delta z = 1$. Our fitting function (Equation ~\ref{eq:sfrd}) places this peak at $z \approx 1.85$, which is plausible but still uncertain. Uncertainties in the faint-end slope of the IRLF and in extinction corrections for the UVLF still dominate at this peak era of cosmic star formation. Although evidence seems to point clearly to a steady increase in the SFRD from $z = 8$ to $z \approx 2$, our direct knowledge of dust-obscured star formation at these redshifts is, for the most part, limited to the rarest and most ultraluminous galaxies, leaving considerable uncertainty about how much SFRD we may be missing in the UV census of that early phase of galaxy evolution. At $z > 4$, our galaxy surveys have been strongly biased toward UV-bright galaxies, and may underestimate both SFRDs and SMDs. Even for UV-selected galaxies, the measurements at $z \geq 8$ are very new and likely uncertain, unsupported by spectroscopic confirmation to date. In addition to measuring redshifts, spectroscopy from the {JWST} will help clarify basic issues about nebular line emission and the degree to which it has affected the photometric analyses that have been carried out to date. Painstaking though all this vast effort has been, it does miss a crucial point. It says little about the inner workings of galaxies, i.e., their ``metabolism'' and the basic process of ingestion (gas infall and cooling), digestion (star formation), and excretion (outflows). Ultimately, it also says little about the mapping from dark matter halos to their baryonic components. Its roots are in optical-IR astronomy, statistics, stellar populations, and phenomenology, rather than in the physics of the ISM, self-regulated accretion and star formation, stellar feedback, and SN-driven galactic winds. It provides a benchmark against which to compare semi-analytic modeling and hydrodynamical simulations of galaxy formation, but it offers little guidance in identifying the smaller-scale basic mechanisms that determine the rate of conversion of gas into stars and lead to the grandiose events in the history of the Universe described in this review. A variety of physical processes are thought to shape the observed distribution of galaxy properties, ranging from those responsible for galaxy growth (e.g., star formation and galaxy merging) to those that regulate such growth (e.g., energetic feedback from SNe, AGN, and the UV radiation background). However, many of these processes likely depend primarily on the mass of a galaxy's dark matter halo. Relating the stellar masses and SFRs of galaxies to the masses and assembly histories of their parent halos is a crucial piece of the galaxy formation and evolution puzzle. With the accumulation of data from large surveys and from cosmological numerical simulations, several statistical methods have been developed over the past decade to link the properties of galaxies to the underlying dark matter structures (e.g. Berlind \& Weinberg 2002, Yang \etal 2003, Vale \& Ostriker 2004). One of them, the ``abundance matching" technique, assumes in its simplest form a unique and monotonic relation between galaxy light and halo mass, and it reproduces galaxy clustering as a function of luminosity over a wide range in redshift (e.g. Conroy \etal 2006, Guo \etal 2010, Moster \etal 2010). Modern versions of this approach (Moster \etal 2013, Behroozi \etal 2013) have shown that {\it a)} halos of mass $\sim 10^{12}\,\msun$ are the most efficient at forming stars at every epoch, with baryon conversion efficiencies of 20-40\% that fall rapidly at both higher and lower masses; {\it b)} in halos similar to that of the Milky Way, approximately half of the central stellar mass is assembled after redshift 0.7; and {\it c)} in low-mass halos, the accretion of satellites contributes little to the assembly of their central galaxies, whereas in massive halos more than half of the central stellar mass is formed ``ex-situ." These studies represent promising advances, albeit with serious potential shortcomings (e.g. Guo \& White 2013, Zentner \etal 2013). The assumption of a monotonic relation between stellar mass and the mass of the host halo is likely incorrect in detail, and it predict only numerically converged properties on scales that are well resolved in simulations. The matching procedure requires minimal assumptions and avoids an explicit treatment of the physics of galaxy formation. As such it provides relatively little new insight into this physics. In the version of this technique by Behroozi \etal (2013), for example, the cosmic SFH is reproduced by construction. As of this writing, a solid interpretation of the cosmic SFH from first principles is still missing (for a recent review, see Mac Low 2013). Generically, one expects that star formation may be limited at early times by the build-up of dark matter halos and quenched at low redshift as densities decline from Hubble expansion to the point where gas cooling is inhibited. These two regimes could then lead to a peak in the SFH at intermediate redshifts (Hernquist \& Springel 2003). A decade ago, hydrodynamical simulations predicted that the peak in star-formation activity should occur at a much higher redshift, $z\gta 5$, than is actually observed (Springel \& Hernquist 2003, Nagamine \etal 2004). Theoretical modeling has been unable to correctly forecast the evolution of the SFRD because of the large range of galaxy masses that contribute significantly to cosmic star formation and the difficulty in following the feedback of energy into the ISM and circumgalactic medium from stellar radiation, SN explosions, and accreting massive black holes. Gas cooling in an expanding Universe is an intrinsically unstable process because cooling acts to increase the density of the gas, which in turn increases the cooling rate. Systems collapsing at low redshift have low mean densities and long cooling times, whereas systems collapsing at higher redshifts have higher mean densities and cool catastrophically. Without feedback processes that transfer energy to the ISM and reheat it, one is faced with the classical overcooling problem -- the unphysical cooling of hot gas in the poorly resolved inner regions of galaxies -- and with the consequent overproduction of stars at early times. And yet, a completely satisfactory treatment of feedback in hydrodynamical simulations that capture large cosmological volumes remains elusive, as these mechanisms operate on scales too small to be resolved and must therefore be incorporated via ad-hoc recipes that are too simplistic to capture the complex subgrid physics involved (e.g., Schaye \etal 2010). In-depth knowledge of the mechanisms responsible for suppressing star formation in small halos (e.g., Governato \etal 2010, Krumholz \& Dekel 2012, Kuhlen \etal 2012), more powerful supercomputers, better algorithms as well as more robust numerical implementations of stellar feedback (e.g., Agertz \etal 2013) all now appear as crucial prerequisites for predicting more realistic SFHs. Newer and deeper observations from the ground and space should improve our measurements of the galaxy population and its integrated properties, especially at and beyond the current redshift frontier where data remains sparse. It seems likely, however, that the most important contribution of new surveys and better modeling will be toward a detailed understanding of the physics of galaxy evolution, not simply its demographics. | 14 | 3 | 1403.0007 |
1403 | 1403.7117_arXiv.txt | We report ALMA and SMA observations of the luminous infrared merger NGC 3256, the most luminous galaxy within $z=0.01$. Our data show that both of the two merger nuclei separated by 5\arcsec\ (0.8 kpc) on the sky have a compact concentration of molecular gas. We identify them as nuclear disks with molecular gas surface densities over $10^3$ \Msol\persquarepc\ and determine that while one at the northern nucleus is nearly face-on the other at the southern nucleus is almost edge-on. The northern nucleus is more massive and has molecular arcs and spiral arms around. The high-velocity molecular gas previously found in the system is resolved to two components. They are two molecular outflows associated with each of the two nuclei. The molecular outflow from the northern nuclear disk is part of a starburst-driven superwind seen nearly pole on. It has a maximum velocity greater than 750 \kms\ and its mass outflow rate is estimated to be $\geq 60$ \Msol\peryr\ for a conversion factor $N_{\rm H_2}/I_{\rm CO(1-0)}$ of $1\times 10^{20}$ \unitofX. The outflow from the southern nucleus is a highly collimated bipolar molecular jet seen nearly edge-on. Its line-of-sight velocity increases with distance out to 300 pc from the southern nucleus. The maximum de-projected velocity is on the order of 2000 \kms\ for the estimated inclination and should exceed 1000 \kms\ even allowing for its uncertainty. The mass outflow rate is estimated to be $>50$ \Msol\peryr\ for this outflow. There are possible signs that this southern outflow has been driven by a bipolar radio jet from an AGN that became inactive very recently. The sum of these outflow rates, although subject to the uncertainty in the molecular mass estimate, either exceeds or compares to the total star formation rate in NGC 3256. The feedback from nuclear activities in the form of molecular outflows is therefore significant in the gas consumption budget, and hence evolution, of this luminous infrared galaxy. | \label{s.introduction} NGC 3256 is an infrared-luminous merger with a bolometric luminosity of $L_{\rm bol}=4\times 10^{11} L_{\odot}$ ($D = 35$ Mpc, see Table \ref{t.4418param} for other parameters). Its two nuclei with a projected separation of 5\arcsec\ = 850 pc \citep{Zenner93,Norris95} and two long tidal tails of stars and \ion{H}{1} gas \citep{English03} indicate that the system is in a late stage of merging between two disk galaxies \citep{Toomre77}. NGC 3256 belongs to the sequence of `most luminous galaxies within their distance ranges', which are, beyond the local group, NGC 253, M82, NGC 1068, NGC 3256, Arp 299, and Arp 220 in the catalogue of \citet{Sanders03}. It is therefore among the best targets to explore luminosity-related phenomena in local galaxies, although its location at Dec. = $-43\arcdeg$ impeded its studies compared to other galaxies in the sequence. \citet[hereafter SHP06]{Sakamoto06} made the first interferometric imaging of a CO line emission in NGC 3256 soon after the commissioning of the Submillimeter Array (SMA) and discovered wide CO line wings underlying the much brighter narrow component in previous observations \citep[e.g.,][]{Sargent89,Aalto02}. The wing CO emission was attributed to a molecular outflow from the face-on merger. The detection of a galactic molecular outflow from faint and wide CO line wings became possible at that time owing in part to the new wide-band capabilities of the SMA. Many extragalactic molecular outflows have been detected since then through broad CO line wings caught with wide-band spectrometers \citep[e.g.,][]{Feruglio10,Chung11,Alatalo11}\footnote{Detection of molecular outflows from broad OH lines dates back much further \citep[and references therein]{Baan89}. Galactic outflows of cold molecular gas have been also found from off-plane molecular gas of edge-on galaxies [e.g., \citet{Nakai87} toward M82 and \citet{Turner85}, \citet{GarciaBurillo00}, and \citet{Bolatto13a} toward NGC 253] and from blueshifted molecular absorption lines against nuclear continuum \citep[e.g.,][for Arp 220 and Mrk 231]{Baan89,Sakamoto09,Fisher10}. All galaxies mentioned here belong to the above-mentioned elite sequence of luminous nearby galaxies.} Such molecular outflows coexist with outflows of ionized and atomic gas and are expected to have significant impact on the luminosity-generation activities in galaxies and the evolution of galaxies themselves \citep[for reviews]{Veilleux05,Carilli13}. We have used the new Atacama Large Millimeter/sub-millimeter Array (ALMA) in its first open-use (Cycle 0) to further study NGC 3256. We aimed at the structure and properties of the molecular gas around the luminous merger nuclei including the high-velocity molecular gas. Although the broad CO wings had been confirmed and found to be even broader in the ALMA commissioning and science verification data \citep{Sakamoto13a} its structure was still largely unconstrained. We therefore observed the galaxy in the 3 and 0.8 mm bands in ALMA Cycle 0 and also made supplemental 1.3 mm observations with the SMA. These new observations provide much higher spatial resolution than before for the circumnuclear molecular gas, up to about 1\arcsec\ for CO(1--0), 0\farcs8 for CO(2--1), and 0\farcs5 for CO(3--2). We also obtained high-resolution high-sensitivity data of CN(1--0), \propyne(6--5), \thirteenCO(2--1), \HCOplus(4--3), and 3 and 0.8 mm continuum. In this paper, we report these new observations and give an overall account of the spatial and kinematical structure of the molecular gas in the center of NGC 3256. We found the high-velocity gas to be two bipolar molecular outflows from the two nuclei and that the two outflows have distinctively different properties from each other. We describe our observations and data reduction in Section \ref{s.obs} and present our observational results in Section \ref{s.result}. We use the data in Section \ref{s.configuration} to constrain the merger configuration, which is critical to interpret the observed gas motion. In Section \ref{s.twoOutflows} we present our two-outflow model for the observed velocity structure and gas distribution. The one from the southern nucleus has remarkable properties in its velocity field, high velocity, high collimation, and large energy. We discuss its driving mechanism in Section \ref{s.Snucleus}. Section \ref{s.conclusions} compares our findings in NGC 3256 with similar objects and phenomena in galaxies and then summarizes our conclusions. | \label{s.conclusions} We have reported our ALMA and SMA observations of molecular line and continuum emission in the center of NGC 3256. We constrained the configuration of the two merger nuclei and their nuclear molecular disks much better than before and resolved for the first time the high-velocity molecular gas in the merger into two molecular outflows from the two nuclei. We have suggested the southern molecular outflow from NGC 3256S to be driven by an AGN bipolar jet. If confirmed, it joins a small group of outflows that share the same driving mechanism and have been imaged in molecular line(s). They include the molecular outflows in M51 \citep{Matsushita04}, NGC 1266 \citep{Alatalo11}, and NGC 1433 \citep{Combes13}. Compared with these outflows, the bipolar molecular jet of NGC 3256S is better collimated and more energetic for a common \Xco. This may be because the AGN radio `jets' in the other galaxies are wider radio plumes. Mainly because of the large outflow velocity, the kinetic luminosity of the southern outflow approaches that of local ultraluminous infrared galaxies and quasar hosts observed by \citet{Cicone14}, who obtained outflow kinetic luminosities on the orders of $10^{36}$--$10^{37}$ W with a conversion factor 3 times lower than ours. The large maximum velocity of the southern outflow is also comparable to or larger than those in their survey but this is probably because ours is helped much by the high ALMA sensitivity and the proximity of NGC 3256. The overall significance of AGN-driven, jet-entrained molecular outflows is an open question. AGN time variability similar to the one we suggested for NGC 3256S may reduce the apparent AGN contribution to galactic molecular outflows. Regarding jet-entrained outflows, on one hand, radio jets have been found only in minority of AGNs. For instance, \citet{Ho01} found ``linear'' structures of radio continuum in 14/52 = 27\% of optically selected, nearby Seyfert galaxies. On the other hand, the parameters of our southern outflow imply that a jet-entrained outflow can be more powerful and efficient than other outflows when normalized by the source bolometric luminosity. It is possible therefore that the small number and/or short lifetime of the outflows driven by AGN radio jets are offset to some extent by their efficiencies and luminosities. The two molecular outflows in NGC 3256 are excellent targets for such assessment because we can simultaneously study properties and driving mechanisms of two powerful molecular outflows of different natures. Our observations have added two similarities between NGC 3256 and Arp 220 in addition to both being late stage mergers with large infrared luminosities. One is the presence of outflows from both of the two merger nuclei; for Arp 220 blueshifted molecular line absorption indicative of outflow has been detected toward both nuclei \citep{Sakamoto09}. The other is that the two merger nuclei with less than 1 kpc projected separation still retain their nuclear gas disks with misaligned rotational axes; for Arp 220 this was first imaged by \citet{Sakamoto99}. Our submillimeter observations also revealed a clear difference between the two mergers. Namely, the nuclei of NGC 3256 are less obscured than the Arp 220 nuclei in terms of gas and dust column density averaged at 100 pc scale. This is most clearly seen in the submillimeter continuum emission whose opacity due to dust is almost unity at 860 \micron\ toward the nuclei of Arp 220 but about two orders of magnitude lower toward the nuclei of NGC 3256. In order for NGC 3256 to evolve into Arp 220, therefore, significant gas accretion is needed to the nuclei despite the ongoing strong molecular outflows that would deplete the gas in the nuclei in Myrs. Such evolution may indeed occur because Arp 220 is probably more advanced as a merger than NGC 3256 judging from their nuclear separations. NGC 3256 may become more luminous in that process, perhaps as luminous as Arp 220, because there is a statistical trend for larger nuclear obscuration (i.e., more gas funneling to the nuclei) and larger total luminosities in more advanced mergers \citep{Haan11,Stierwalt13}. Further studies on NGC 3256 are warranted also for the purpose of tracing the late evolutionary path of a merger that is plausibly about to become ultraluminous. Finally we reemphasize our caution on \Xco\ in particular for the high-velocity molecular outflows. The large line widths of the outflow gas reduce the CO column density per line width and hence may well result in optically thin CO emission. The conversion factor for that case is $\Xtwenty \sim 0.1$. Such a low conversion factor for optically-thin CO has been adopted, for example, for the molecular outflow in NGC 1266 on the basis of multi-line CO excitation analysis \citep{Alatalo11}. The outflows in NGC 3256 may have a similar situation and \Xco. Alternatively, the outflowing gas may consist of an ensemble of optically-thick (in CO) clouds that spread in a wide velocity range. In its partial support is our detection of CN(1--0) lines, with likely enhancement relative to CO(1--0), in the high-velocity gas (\S\ref{s.result.line.hv.spectra}). Although CN may be subthermally excited, the detection of a line with a $10^6$ \percubiccm\ critical density implies gas clumping for dense gas to exist in the high velocity outflows. Even if individual clumps are not virialized as assumed for the standard \Xco, the conversion factor for optically thick clumps will be larger than that for optically thin CO (and lower that that for virialized CO-thick clouds). Similar clumping and presence of dense gas in a galactic molecular outflow have been deduced for Mrk 231 by \citet{Aalto12} from their detection of broad line wings in HCN, \HCOplus, and HNC lines. Because most outflow parameters in Table~\ref{t.4418measured.param} depend on \Xtwenty, followup studies on the physical and chemical properties of the high-velocity gas are highly desired. Our primary findings are: 1. Each of the two merger nuclei has its own nuclear disk where molecular line and continuum emission peak. The northern nuclear disk is nearly face-on ($i$ \about30\arcdeg), has a \about200 pc characteristic radius, and clearly rotates around the northern nucleus. The southern nucleus has a more compact emission peak and a linear structure extending \about200 pc on either side. It is deduced to be a nearly edge-on nuclear disk rotating around the southern nucleus. The mean molecular gas surface densities of both nuclei is about 3$X_{20} \times10^4$ \Msol\persquarepc\ at 240 pc resolution, where $X_{20}$ is the CO-to-\HH\ conversion factor in the unit of $10^{20}$ \unitofX. The peak gas surface density is $6 X_{20} \times10^4$ \Msol\persquarepc\ at the southern nucleus at 80 pc resolution. 2. The high velocity molecular gas previously found at the center of the merger is resolved to two molecular outflows associated with the two nuclei. We detected not only CO but also CN lines with enhancement in these outflows. The CN detection in a galactic outflow is for the first time to our knowledge. The total molecular outflow rate of the two outflows is on the same order of the total star formation rate in NGC 3256. 3. The molecular outflow from the northern nuclear disk is a bipolar flow with a wide opening angle and a nearly pole-on viewing angle. It has de-projected outflow velocities up to 750 \kms\ at $\gtrsim$4$\sigma$ and an outflow time scale (crossing time) of 1 Myr. Its molecular gas mass is $6 X_{20} \times10^7$ \Msol, mass outflow rate $60 X_{20}$ \Msol\peryr, and kinetic luminosity on the order of $4 X_{20} \times 10^8$ \Lsun. The last three are for the gas at de-projected velocities above 260 \kms. At the current rate the outflow would deplete molecular gas in the northern nuclear disk in 3 Myr if the same conversion factor applies to the nuclear disk and the outflow. Most of the outflow/superwind signatures found so far at other wavelengths in NGC 3256 must be from this outflow. 4. The molecular outflow from NGC 3256S is a well collimated bipolar jet with a \about$20\arcdeg$ opening angle and is nearly edge on. It has a de-projected maximum velocity 2600 \kms\ for a favored inclination angle 80\arcdeg\ or 1300 \kms\ for $i=70\arcdeg$. The line-of-sight outflow velocity increases with distance up to 300 pc from the nucleus. This molecular jet has a 0.5 Myr crossing time, a mass of $2.5 X_{20} \times10^7$ \Msol, a mass outflow rate $50 X_{20}$ \Msol\peryr, and a kinetic luminosity on the order of $90 X_{20} \times 10^8$ \Lsun\ for $i=80\arcdeg$. These are for gas at projected velocities above 220 \kms\ and the lower velocity gas in the outflow may be an order of magnitude larger in mass. The gas depletion time for the central 80 pc is \about0.6 Myr under the same assumption about the conversion factor as above and ignoring the lower velocity flow. 5. The northern outflow is a starburst driven superwind in all likelihood. The southern outflow is most likely entrained by a radio jet from a weak or recently dimmed AGN in the southern nucleus. Pieces of evidence for the latter outflow driver are the large differences in the outflow parameters from the northern superwind, off-nuclear broad \Halpha\ lines in NGC 3256, and a pair of radio spurs from the southern nucleus that matches in shape the southern molecular bipolar jet. 6. Continuum spectral indexes are negative at 3 mm and positive at 0.86 mm for both nuclei. The index is lower, in particular at 0.86 mm, for the southern nucleus, suggesting significant synchrotron and/or free-free emission even at 860 \micron. Neither nucleus has a bright ($\Tb > 10$ K) dust continuum core of several 10 pc size at 860 \micron\ such as those found in Arp 220 and NGC 4418. This disfavors presence of a highly Compton-thick and currently luminous AGN in the nuclei of NGC 3256. The new observations presented in this paper contain more information than we could fit in a single paper. Further analysis will be reported elsewhere. \vspace{5mm} | 14 | 3 | 1403.7117 |
1403 | 1403.4114_arXiv.txt | At low energies nucleon-nucleon interactions are resonant and therefore supernova matter at subnuclear densities has many similarities to atomic gases with interactions dominated by a Feshbach resonance. We calculate the rates of neutrino processes involving nucleon-nucleon collisions and show that these are enhanced in mixtures of neutrons and protons at subnuclear densities due to the large scattering lengths. As a result, the rate for neutrino pair bremsstrahlung and absorption is significantly larger below $10^{13}$\,g\,cm$^{-3}$ compared to rates used in supernova simulations. | 14 | 3 | 1403.4114 |
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