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1206 | 1206.6105_arXiv.txt | We provide evidence that the obliquities of stars with close-in giant planets were initially nearly random, and that the low obliquities that are often observed are a consequence of star-planet tidal interactions. The evidence is based on 14 new measurements of the Rossiter-McLaughlin effect (for the systems HAT-P-6, HAT-P-7, HAT-P-16, HAT-P-24, HAT-P-32, HAT-P-34, WASP-12, WASP-16, WASP-18, WASP-19, WASP-26, WASP-31, Gl\,436, and Kepler-8), as well as a critical review of previous observations. The low-obliquity (well-aligned) systems are those for which the expected tidal timescale is short, and likewise the high-obliquity (misaligned and retrograde) systems are those for which the expected timescale is long. At face value, this finding indicates that the origin of hot Jupiters involves dynamical interactions like planet-planet interactions or the Kozai effect that tilt their orbits, rather than inspiraling due to interaction with a protoplanetary disk. We discuss the status of this hypothesis and the observations that are needed for a more definitive conclusion. | \label{sec:intro} \setcounter{footnote}{0} Exoplanetary science has been full of surprises. One of the biggest surprises emerged at the dawn of this field: the existence of ``hot Jupiters'' having orbital distances much smaller than an astronomical unit (AU). It is thought that giant planets can only form at distances of several AU from their host stars, where the environment is cooler and solid particles are more abundant, facilitating the growth of rocky cores that can then attract gaseous envelopes from the protoplanetary disk. Different mechanisms have been proposed which might transport giant planets from their presumed birthplaces inward to where we find them. Among the differences between the proposed mechanisms is that some of them would alter the planet's orbital orientation and thereby change the relative orientation between the stellar and orbital spin \citep{nagasawa2008,fabrycky2007,naoz2011}, while others would preserve the relative orientation \citep{lin1996}, or even reduce any primordial misalignment \citep{cresswell2007}. For this reason, measuring the stellar obliquity---the angle between stellar and orbital axes---has attracted attention as a possible means of distinguishing between different theories for the origin of hot Jupiters. The stellar obliquity is an elusive parameter because the stellar surface needs to be at least partially resolved by the observer. If the system exhibits eclipses or transits, then it is possible to detect anomalies in the spectral absorption lines of the eclipsed star, which have their origin in the partial blockage of the rotating photosphere. The precise manifestation of the rotation anomaly depends on the angle between the projections of the stellar rotation axis and orbital axis of the occulting companion. Credit for the first definitive measurements has been apportioned to \cite{rossiter1924} and \cite{mclaughlin1924}, after whom the effect is now named. \cite{queloz2000} were the first to measure the Rossiter-McLaughlin (RM) effect for a planet-hosting star, finding a low obliquity. Since then many systems have been studied, including those with misaligned stars \citep{hebrard2008}, retrograde orbits \citep{winn2009,triaud2010}, a hot Neptune \citep{winn2010c,hirano2011} and even a circumbinary planet \citep{winn2011b}. \cite{winn2010} found a possible pattern in the hot-Jupiter data, namely, host stars hotter than $T_{\rm eff} \approx 6250$~K tend to have high obliquities, while cooler stars have low obliquities. \cite{schlaufman2010}, using a different method, also found a preponderance of high obliquities among hot stars. \cite{winn2010} speculated that this pattern was due to tidal interactions. Specifically, it was hypothesized that hot Jupiters are transported inward by processes that perturb orbital inclinations and lead to a very broad range of obliquities. Cool stars ultimately come into alignment with the orbits because they have higher rates of tidal dissipation, due to their thick convective envelopes. Hot stars, in contrast, lack thick convective envelopes and are unable to reorient completely on Gyr timescales. If this interpretation is correct, then not only should further measurements be consistent with the hot/cool pattern, but also the degree of alignment should be found to correlate with the orbital period and planet-to-star mass ratio, parameters which also strongly affect the rate of tidal dissipation. Since the study by \cite{winn2010} the number of systems with RM measurements has nearly doubled. Here we present the results of an additional 14 observations, as well as a critical review of other published measurements (including thorough re-analyses in three cases). We can now attempt a comparison between the measured obliquities and the theoretical distribution of obliquities that one would expect if tides were an important factor. We also refer the reader to \cite{hansen2012}, who recently performed a comparison with a similar motivation, without the benefit of the new RM measurements presented here, but using more sophisticated theoretical models and drawing qualitatively similar conclusions. This paper has two main parts. The first part is observational. We describe our new observations of the RM effect (Section~\ref{sec:obs}), our analysis method (Section~\ref{sec:analysis}), and the details of individual systems analyzed here (Section~\ref{sec:systems}). The second part (Section~\ref{sec:discussion}) considers the distribution of obliquities, seeks evidence for the expected signatures of tidal effects, and considers the implications for the origin of hot Jupiters. We summarize our results in Section~\ref{sec:summary}. | \label{sec:discussion} We now use the measurements of projected obliquities obtained in the last section together with other measurements taken from the literature to learn more about the evolution of these systems.\footnote{We present here all measurements with the sign convection form \cite{ohta2005} and use the symbol $\lambda$. Often researchers use the symbol $\beta$ and the sign convention from \cite{hosokawa1953}, where $\lambda = -\beta$.} In \ref{sec:remarks} we first remark on some specific cases from the literature. In Section \ref{sec:tides} we revisit the evidence for an correlation between the degree of alignment and the effective stellar temperature discovered by \cite{winn2010}. We then discuss further evidence of tidal interaction, based on correlations between the planet-to-star mass ratio and $\lambda$, and the dependence of $\lambda$ on the orbital distance between the two bodies. In Section \ref{sec:tides} we sort the systems according to a theoretical tidal timescale, as a test of whether tides have been important in altering the obliquities of these systems. We also discuss the implications for the origin of hot Jupiters and the strengths and weaknesses of our interpretation. \subsection{Remarks on the previous literature} \label{sec:remarks} \paragraph{CoRoT--1} The RM effect was measured by \cite{pont2010}. We do not include this measurement in our subsequent discussion because of the low SNR of the detection and because it was only possible to obtain a few out-of-transit observations. The authors caution that systematic uncertainties could cause the actual errors in the measurement of the projected obliquity ($\lambda=77\pm11^{\circ}$) to be larger than the statistical uncertainty. We note that if the value for $\lambda$ from this study were taken at face value, it would constitute an exception to the pattern presented below. For this reason it would be an important system to reobserve. It is a challenging target because of the faint and slowly-rotating host star. \paragraph{CoRoT--11} \cite{gandolfi2010} obtained RVs during a planetary transit in this system. As only the first half of the transit was observed, and only with a low SNR, they could only conclude that the orbit was prograde. In that sense this is a similar case to HAT-P-16 presented here. We decided to exclude CoRoT-11 from subsequent discussion as we did with our result on HAT-P-16. \paragraph{CoRoT--19} This system has an F9V star and a Jupiter-mass planet on a $3.9$-day orbit. The RM effect was measured by \cite{guenther2012}. They found $\lambda=-52^{+27}_{-22}\,^{\circ}$, a misalignment between stellar and orbital axes. However the RM effect was only detected at the 2.3$\sigma$ level, and no post-egress data were obtained. We omit this measurement in the subsequent discussion. \paragraph{KOI--13} \cite{szabo2011} detected a slight asymmetry in the transit light curve and attributed the asymmetry to a misalignment of the planet's orbit relative to the stellar rotation. The host star is a fast rotator, leading to a lower surface gravity and surface brightness around the equator compared to the poles. By modeling this effect, \cite{barnes2011} found $|\lambda|$ to be either $24\pm4^{\circ}$ or $156\pm4^{\circ}$. Either choice represents a substantial misalignment and would lead to similar conclusions in the subsequent discussion. For simplicity of presentation in the plots to follow, we arbitrarily adopt the lower value. \cite{barnes2011} also calculated the stellar inclination along the line of sight, which we do not use here, since this information is not available for most of the other systems. We adopt the mass for the secondary estimated by \cite{mislis2012}, which is also consistent with work by \cite{shporer2011c}. Both estimates were based on the photometric orbit. We further adopt the age estimate of ($710^{+180}_{-150}$~Myr) by \cite{szabo2011}. \paragraph{WASP-23} We omit the measurement by \citep{triaud2011b}. Because of a low impact parameter, the only conclusion that could be drawn is that orbit is prograde. \paragraph{XO--2} We omit the measurement by \citep{narita2011}. They found $\lambda=10\pm72^{\circ}$, a prograde orbit but with a very large uncertainty, similar to our result for HAT-P-16. \subsection{Relevant system properties} \label{sec:tides} \begin{figure*} \begin{center} \includegraphics{teff.pdf} \caption {\label{fig:proj_obli} {\bf Projected obliquities and projected stellar rotation speeds as a function of the stellar effective temperature.} {\it Upper panel:} Measurements of projected obliquities as a function of the effective temperature of the host star. Stars which have temperatures higher then $6250$~K are shown with red filled symbols. Blue open symbols show stars with temperatures lower then $6250$~K. Stars which measured effective temperature include $6250$~K in their 1-$\sigma$ interval are shown by split symbols. Systems which harbor planets with mass $<0.2$~$M_{\rm Jup}$ or have an orbital period more then $7$~days are marked by a black filled circle with a ring. {\it Lower panel:} Projected stellar rotation speeds $v \sin i_{\star}$ of the stars in our sample. In addition $v \sin i_{\star}$ measurements of stars in the catalogue by \cite{valenti2005} are shown as small dots.} \end{center} \end{figure*} \paragraph{Effective temperature} \cite{winn2010} noted that for hot Jupiters, nonzero values of $\lambda$ tended to be associated with hot stars. For effective temperatures $\gtrsim6250$~K the obliquities have a broad distribution, while for lower temperatures the measurements are consistent with low obliquities. The only exceptions were two systems with significantly longer orbital periods than the rest. \cite{schlaufman2010} independently found that hot stars tended to be misaligned with the orbits of hot Jupiters by comparing the measured value of $v\sin i_{\star}$ with the expected value for $v$, for a star of the given mass and age. While this approach has the virtue of requiring less intensive observations, it does rely heavily on accurate measurement of $v\sin i_{\star}$, which is problematic for slowly-rotating stars, and assumes that $v$ was not affected by any tidal influence of the close-in planet. \begin{figure*} \begin{center} \includegraphics{mass_ratio.pdf} \caption {\label{fig:mass_ratio} {\bf Projected obliquities as function of the planetary to stellar mass ratio}. The same symbols as in Figure~\ref{fig:proj_obli} are used. With larger mass of the companion the degree of misalignment decreases, while for some of the more massive planets a small but significant obliquity is detected. These planets orbit stars with radiative envelopes, reducing the effectiveness with which tidal energy can be dissipated. Names of systems with particular small or large mass ratios are indicated.} \end{center} \end{figure*} With the new measurements presented in this paper and with measurements by others over the last two years, the number of systems with measured projected obliquities is now up to 53. This is more than twice the number that was available to \cite{winn2010}. Table~\ref{tab:proj_obli} shows the measured projected obliquities of all systems used in this study.\footnote{See also {\tt exoplanets.org, exoplanet.eu, www.astro. keele.ac.uk/jkt/tepcat}, and {\tt www.aip.de/People/rheller} for listing of obliquity measurements.} The increased number in measurements enables a stringent test of the proposed pattern, as well as a more in depth analysis and comparison to other system parameters. It should be noted that in almost all cases, the measurements are of the projected obliquity, and not the true obliquity. For true obliquities smaller then $90^{\circ}$ the projected obliquity is usually {\it smaller} than the true obliquity, while for true obliquities $>$$90^{\circ}$ the projected obliquity is usually {\it larger} than the true obliquity. This factor complicates any detailed comparison between the measurements and the theoretical expectations. For simplicity we have chosen to work with projected obliquities, rather than attempting any deprojection scheme \citep{fabrycky2009,morton2011}. The upper panel of Figure~\ref{fig:proj_obli} shows the projected obliquities plotted as a function of the effective temperature of the host star. Apparently, for these systems, the trend observed by \cite{winn2010} still holds. There are three apparent exceptions to the rule: HAT-P-11, WASP-8, and HD\,80606. These systems are special in other ways too, by virtue of having either an unusually low planet mass or an unusually long orbital period. They represent three of the four systems for which the orbital period is greater than 7~days or the planet has a mass lower than $0.2$~$M_{\rm Jup}$. In this sense they least resemble the typical ``hot Jupiter''. We will discuss these important cases in the following paragraphs. The explanation for the relationship between $T_{\rm eff}$ and $\lambda$ could fall into one of two categories: (i) the formation and evolution of hot Jupiters is different for hot stars than for cool stars, which for some reason results in higher obliquities in the hot stars. (ii) The distribution of obliquities is originally broad for both hot stars and cool stars, but they evolve differently depending on a parameter closely associated with $T_{\rm eff}$. \cite{winn2010} suggested that the second scenario is more likely and that the factor associated with temperature is the rate of tidal dissipation due to the tide raised by the planet. The reason for this suspicion was that $T_{\rm eff} \approx 6250$~K is not an arbitrary temperature, but rather represents an approximate boundary over which the internal structure of a main-sequence star changes substantially. Stars hotter than this level have very thin or absent convective envelopes, with the mass of the envelope dropping below about $0.002~M_{\odot}$ at 6250~K \citep{pinsonneault2001}. (For the Sun, the mass of the convective envelope is around $0.02~M_{\odot}$). Independently of this theoretical expectation, there is dramatic empirical evidence for a transition in stellar properties across the 6250~K divide: hotter stars are observed to rotate more rapidly. In the lower panel of Figure~\ref{fig:proj_obli}, we plot the projected rotation speeds of a sample of $\sim1000$ stars from the catalogue by \cite{valenti2005}. The projected rotational speed $v \sin i_{\star}$ increases rapidly around 6250~K. For F0 stars, the rotation speed can approach $200$~km\,s$^{-1}$. It is thought that stellar rotation together with the convection in the envelope create a magnetic field coupling to the ionized stellar wind far beyond the stellar radius and thereby transport angular momentum away from the stellar rotation \citep[see, e.g., ][for further discussion.]{barnes2003} Presumably this magnetic braking is less effective for stars without convective envelopes, leading to the observed rapid increase in stellar rotation speeds towards earlier spectral type. Judging from Figure~\ref{fig:proj_obli}, the transition from low obliquity to high obliquity seems to be linked empirically to this transition from slowly-rotating to rapidly-rotating stars. The presence of a convective envelope is also expected to change the rate of dissipation of the energy in tidal oscillations. Energy contained in tidal bulges is thought to be more effectively dissipated by turbulent eddies in convective envelopes than by any mechanism acting in a radiative envelope. See \cite{zahn2008} for a review on the theory of tidal interactions. \cite{torres2010} and \cite{mazeh2008} review the evidence for tidal interactions in double-star systems, and provide more access points to the literature. If tidal evolution is responsible for the difference in stellar obliquities between cool and hot stars, then there should also be a correlation between the mass ratio of the star and planet and the degree of alignment. The higher the mass of the planet ($m_{\rm p}$) relative to the mass of the star ($M_{\star}$), the faster its tides can align the stellar and orbital angular momentum vectors. Furthermore, one would expect an inverse correlation between the scaled semi-major axis ($R_{\star}/a$) and the obliquity. A more distant companion will raise smaller tides. We will next investigate whether such correlations exist. We will also investigate if the age of the systems is an important factor in setting the degree of alignment in these systems. \paragraph{Mass Ratio} In Figure~\ref{fig:mass_ratio} the measured projected obliquities are plotted as a function of the mass ratio between planet and host star. Higher obliquities are measured for systems in which the mass of the planet is relatively small. This is what would be expected if tides are responsible for the obliquity distribution. Massive planets raise stronger tides. This trend was observed before \cite[e.g.][]{hebrard2011b}, though at that time it was not interpreted in terms of tidal interaction. The interpretation is not clear from this comparison alone, though. Note for example that the most massive planets are found around hot stars, which should have weaker dissipation that counteracts the effect of the more massive planet to at least some degree. In a subsequent section we will attempt to take both these effects (and that of orbital distance) into account at once. \paragraph{Distance dependence} Does the degree of alignment depend on the scaled distance ($a/R_{\star}$)? Figure~\ref{fig:ar} gives a mixed answer to this question. Focusing on systems with cool stars (blue open circles) there seems to be a trend of obliquities as a function of the scaled distance. The data suggest good alignment for all systems with $a/R_{\star}< 10$. Three of the four systems with greater distances have significant misalignments. No such dependence is observed for systems with stars which are close to $6250$~K or hotter. However the range of distances that is spanned by the hot-star systems is very small, less then an order of magnitude. In contrast, the cool-star systems probe nearly two orders of magnitude. Note that two of the close-in, misaligned systems are the systems with the hottest host stars (WASP-33 and KOI-13). This lack of alignment finds an explanation in the tidal hypothesis: despite the tight orbits, the tidal dissipation rate may be relatively low due to thin or non-existent convective layer. \begin{figure*} \begin{center} \includegraphics{ar.pdf} \caption {\label{fig:ar} {\bf Dependence of $\lambda$ on the scaled orbital distance.} The same as Figure~\ref{fig:mass_ratio}, but now the measurements of the projected obliquity are plotted as a function of the semi-major axis divided by the stellar radius. While a trend with distance can be observed for planets around cooler stars, such a trend seems to be absent for planets around hotter stars. This might be due to the small range probed in distance by the obliquity measurements around hot stars. However because tidal forces are week in these stars only for the innermost massive planets (e.g. WASP-12\,b and WASP-18\,b) such a trend would be readily observable. Names of systems with scaled distances greater then $15$, and for misaligned close in hot systems are indicated.} \end{center} \end{figure*} \paragraph{Age} \label{sec:age} Under the tidal hypothesis, older systems should tend to be closer to alignment than younger systems, all else being equal. This is because in older systems, tides have had a longer interval over which to act. Included in ``all else being equal'' is the underlying assumption that the arrival time of the hot Jupiter to its close-in orbit is the same in all systems. \cite{triaud2011} presented empirical evidence that the degree of misalignment depends chiefly on the age of the system. He found that all systems in his sample with ages greater than $2.5$~Gyr are aligned (see his Figure~2). His sample consisted only of those stars with an estimated mass greater than $1.2$~$M_{\odot}$, since it is harder to determine a reliable age for lower-mass systems. Stars with a mass of $\gsim$1.2~$M_{\odot}$ develop a significant convective envelope during their main sequence lifetime, even if they were too hot to have a significant convective envelope on the ``zero age'' main sequence. In Figure~\ref{fig:age} we plot the projected obliquities as function of stellar age for stars with $M_{\star}>1.2 M_{\odot}$. And indeed all the aligned and older systems are cool enough to have a significant convective envelope. Figure~\ref{fig:age} represents, therefore, a similar pattern as seen in Figure~\ref{fig:proj_obli} with a slight shift in the variable and for a subset of systems (only stars with $M_{\star}>1.2 M_{\odot}$). It seems as though the development of a convective envelope with age, rather than the age itself, might be driving the degree of alignment. \begin{figure} \begin{center} \includegraphics[width=8.cm]{age.pdf} \caption {\label{fig:age} {\bf Projected obliquity plotted as function of age for stars with $M_{\star}>1.2M_{\odot}$}. This is a similar plot to the one presented by \cite{triaud2011}. Same symbols used as in Figure~\ref{fig:proj_obli}. Systems which are older than $\sim3$~Gyr are cool enough to develop a convective envelope. This plot is therefore a relative to Figure~\ref{fig:proj_obli}.} \end{center} \end{figure} \subsection{Tidal timescale} \label{sec:timescale} As we have seen in the last section that the degree of alignment is correlated with stellar temperature, the mass ratio, and possibly the orbital distance. We now try in this section to establish a single quantitative relationship between the degree of alignment and those parameters. Ideally we could calculate a theoretical alignment timescale for each system, and compare that timescale to the estimated age of the system. We could then check if systems with a relatively short timescale (fast alignment) tend to have low obliquities, and systems with timescales comparable to the lifetime of the system (or larger) tend to have high obliquities. \begin{figure*} \begin{center} \includegraphics{tau.pdf} \caption {\label{fig:tau} {\bf Measured projected obliquity as function of the alignment timescale calibrated from binary studies.} The same symbols as in Figure~\ref{fig:proj_obli} are used. This time the projected obliquities are shown as a function of the characteristic timescale needed to align the stellar and orbital axes. We used two different equations to calculate these timescales. One for stars with temperatures lower than $6250$~K for which we assume that tidal dissipation happens due to Eddies in the convective envelope and one for hotter stars for which we assume that no convective envelope is present and alignment is due to radiative damping. The coefficients for these equations have been calibrated with synchronization timescales in double star systems. Note that both timescales have been divided by $5\cdot 10^{9}$. } \end{center} \end{figure*} Calculating timescales needed to synchronize and align stellar rotation is a complex task. Apart from the parameters mentioned above, there are other parameters that would influence the time needed for alignment. For example the total amount of angular momentum stored in the stellar rotation, and the driving frequency of the tidal force (i.e.\ twice the orbital frequency), are expected to be important. In addition, the rate of dissipation is not expected to be constant over Gyr timescales due to the contemporaneous evolution in orbital distance and eccentricity, and due to stellar evolution. Even worse, the specific mechanisms for dissipating tidal energy are not completely understood, neither for stars with radiative envelopes nor for stars with convective envelopes. Nevertheless there are some simple considerations we may employ to obtain approximate timescales for alignment. \begin{enumerate} \item We can use the formulae provided by \cite{zahn1977} for synchronization. The coefficients in these formulae are difficult to derive from theory alone, but they can be calibrated with observations made in binary star systems. By observing the maximum orbital distance within which binary stars are observed to be spin-orbit synchronized, and knowing the lifetime of the stars on the main sequence, the relevant parameters can be estimated. To apply this to our sample two different formulae are needed. One for cool stars which have convective envelopes (CE) and hot stars which have radiative envelopes (RA). Therefore this approach has the virtue of empirical calibration, although the calibration is for star-star interactions rather than planet-star interactions, and the calibration is for spin synchronization rather than reorientation. One complication is that to apply these formulae we have to make a binary decision on whether a star is ``convective'' or ``radiative'' which does not reflect the gradual thinning of the convective envelope with increasing stellar temperature. We choose a temperature of $6250$~K for this boundary. \item Assuming that the alignment timescale due to dissipation in convective envelopes ($\tau_{\rm CE}$) is always shorter than the time needed for alignment by forces in radiative envelopes ($\tau_{\rm RA}$) we can try to derive a simple relationship between the mass contained in the convective envelope and the alignment timescale $\tau$. This would have the advantage that the gradual decrease of mass in the convective envelope can be easily incorporated, but we ignore here any possible additional dissipation mechanism in the radiative envelope which for higher temperatures would become important. In addition it is not obvious why $\tau^{-1}$ should depend linearly on the mass contained in the convective envelope, nor can we be completely confident in our estimate of the convective mass based only on the observable parameters of the stellar photosphere. And of course the convective mass is not really a constant over Gyr timescales. \end{enumerate} \begin{figure*} \begin{center} \includegraphics{tau2.pdf} \caption {\label{fig:tau2} {\bf Measured projected obliquity as a function of the alignment timescale estimated from the mass of the convective envelope divided by age. } Similar to Figure~\ref{fig:tau}. This time however a tidal timescale was calculated which depends next to the mass ratio and the scaled distance not on a calibrated coefficient but on the mass in the convective envelope. Further the age estimates of the systems have been taken into account for this ranking. KOI-13 is not shown in this plot. The mass contained in its convective envelope is small and therefore according to equation~\ref{equ:tidal_time} it does practically not realign.} \end{center} \end{figure*} The simplifications made by either approach should cause us not to expect a perfect and deterministic relationship between our theoretical parameters and the observed obliquities [and we direct the reader to \citet{hansen2012} for a different approach to this comparison]. For the first approach we obtain the following relationships between the system parameters and the convective and radiative timescale for alignment from \cite[][and references therein]{zahn1977}, \begin{equation} \label{equ:ce} \frac{1}{\tau_{\rm CE}} = \frac{1}{10 \cdot 10^{9} {\rm yr}} q^2 \left(\frac{a/R_\star}{40}\right)^{-6} \,{\rm and}\, , \end{equation} \begin{equation} \label{equ:ra} \frac{1}{\tau_{\rm RA}} = \frac{1}{0.25 \cdot 5 \cdot 10^{9} {\rm yr}} q^2(1+q)^{5/6} \left(\frac{a/R_\star}{6}\right)^{-17/2}\,\,\,. \end{equation} Here $q$ is the planet-to-star mass ratio. For stars with convective envelopes, synchronization is observed for binaries out to a scaled distance of $\sim40$ during their main-sequence life, which we set to $10$~Gyr for all 'cool' stars. For 'hot' stars, synchronization in binaries is observed out to $\sim6$ times the scaled radius. Tidal damping is expected to be most efficient during the first quarter of their main sequence life \citep{zahn1977} (which we set here to $5$~Gyr). See e.g. \cite{claret1997} for a more recent comparison between theory and observations, mainly for stars of higher mass. In Figure~\ref{fig:tau} we show the projected obliquity versus the characteristic timescale needed for realignment. For this we used equation~\ref{equ:ce} for $T_{\rm eff} < 6250$~K and equation~\ref{equ:ra} for $T_{\rm eff} \ge 6250$~K. (Both timescales were divided by $5\cdot 10^{9}$ for normalization.) For most of the systems for which rapid alignment is expected, low projected obliquities are observed. For systems where tides are expected to be too slow in aligning and synchronizing stellar rotation, a very broad range of projected obliquities (apparently random) is observed. We now return to the three apparent outliers from Figure~\ref{fig:proj_obli}: HAT-P-11, WASP-8, and HD\,80606. While these are all stars with convective envelopes, the timescales for alignment are very long. This is because the scaled distances are greater than $15$ (Figure~\ref{fig:ar}). In the case of HAT-P-11 there is also the additional penalty from the relatively small planet mass ($q\approx 10^{-4}$). Thus, in this light, these three ``outliers'' are not exceptions; they have high obliquities because the tidal timescales are very long. There is one system which does seem to be an exception: HAT-P-32. The rotation axis of the star nearly lays in the plane of the orbit ($\lambda=85\pm1.5^\circ$), while all other systems with a similar tidal timescale do have projected obliquities consistent with good alignment. Because of the obliquity near $90^{\circ}$ tides couple only weakly to the stellar rotation \citep{lai2012}. The timescale for alignment in this system could be longer than estimated by our simple formula. If this is the only reason for the high obliquity in the HAT-P-32 system then we might expect to find more systems with short alignment timescales and similar obliquities in Figure~\ref{fig:tau}, which do not exist in the current sample. The measured effective temperature in HAT-P-32 is $6207\pm88$~K, and therefore we used Equation~\ref{equ:ce}, applicable to stars with convective envelopes. If we would have assumed an effective temperature of $6250$~K, a value within the $1$-$\sigma$ interval of the measurement, and used Equation~\ref{equ:ra} then we would have obtained a $\tau_{\rm RA}$ of $1.5 \cdot 10^{5} \times 5\cdot 10^{9}$~yr instead of $\tau_{\rm CE} = 1.8\cdot 10^{1} \times 5\cdot 10^{9}$ ~yr. This illustrates the aforementioned weakness of this first approach, that we have to make a binary decision on whether a star is ``convective'' or ``radiative.'' Note also that KOI-13 and XO-3 are ``hot'' stars which have significant misalignments, and yet they are found in between ``cool'' aligned systems. KOI-13 ($T_{\rm eff}=8500$~K) is hottest star in our sample and it is questionable if we can use the same tidal timescale for this system as for the other hot systems which are about $2000$~K cooler (see Figure~\ref{fig:proj_obli}). In the second approach we build upon Equation~\ref{equ:ce}. Now we do not use any empirical calibration. For each planet-hosting star we estimate the mass contained in the convective envelope, and assume that the rate of energy dissipation is proportional to this convective mass, \begin{equation} \label{equ:tidal_time} \frac{1}{\tau} = C \cdot \frac{1}{M_{\rm cz}} q^2 \left(\frac{R}{a}\right)^6 \,\, , \end{equation} where $M_{\rm cz}$ indicates the mass contained in the outer convective zone and $C$ is an unspecified proportionality constant with units g~s$^{-1}$. Our estimate for $M_{\rm cz}$ is based on the measured $T_{\rm eff}$. This ensures a gradual decrease of tidal forces with increasing temperature. To establish the relation between $T_{\rm eff}$ and convective mass we used the EZ-Web tool\footnote{This tool is made available by Richard Townsend under the following url: {\tt http://www.astro.wisc.edu/{\tt\~{}}townsend}} for stars with $T_{\rm eff}<7000~K$, and the data from \cite{pinsonneault2001} for hotter stars. To create Figure~\ref{fig:tau2} we further divided the tidal timescale by the estimated main-sequence age, taking the uncertainty in the age estimate into account. As the ages are not well known, this leads to a substantial uncertainty in the positioning of a system on the horizontal axis. On average, the hot stars are younger than the cool stars. Therefore the main effect of the division by age is a small shift of the hotter stars to the right side of the logarithmic plot. The ordering of the cool stars is not substantially changed, relative to Figure ~\ref{fig:tau}, but there are now a few hot stars with significant obliquities and with similar tidal timescales as some cool aligned systems. The biggest outlier in this respect is HAT-P-7. In summary, despite the shortcomings of our highly simplified approaches to calculating the theoretical tidal timescale, and a few exceptional cases, we do find support for the claim that the obliquities in hot Jupiter systems undergo damping by tidal dissipation. Systems with short tidal timescales are predominantly well-aligned, while systems with longer tidal timescales display an apparently random obliquity distribution. The implication is that the obliquities were once even more broadly distributed than we observe them today. Put differently, the ``primordial'' orbits of hot Jupiters (the orbits that existed shortly after the planets arrived close to the star) once had a very broad range of inclinations relative to the stellar equatorial plane. \subsubsection{Angular momentum problem} As \cite{winn2010} already pointed out, there is a theoretical problem with invoking tides in this context. The angular momentum in the stellar rotation compared to the angular momentum in the orbit (when the planet is close enough to significantly effect the stars rotation via tides) is so large that to synchronize and align the star the planet would surrender so much angular momentum that it would spiral into the star. For nearly all systems in our sample the orbital velocity (at periastron) is larger than stellar rotation velocity. This causes trailing tides, and angular momentum is transported from the orbit towards the stellar rotation, leading to decay of the orbit \citep[e.g.][]{levrard2009}. Yet we see systems which have aligned axes and the planets have evidently survived. To address this problem \cite{winn2010} speculated that only the outer layers of the star synchronize and align with the orbit. In that case a smaller amount of the angular momentum would be transferred and the planetary inspiral would be avoided. It seems doubtful, though, that a separate rotation speed and rotation direction for the envelope relative to the stellar interior could be maintained for billions of years. More recently \cite{lai2012} suggested that the angular momentum problem is not as serious as it might seem. Given the complexities of tidal dynamics, he argued that there is no strong theoretical reason why the timescale for realignment must equal the timescale for synchronization, and indeed he provided a particular theoretical tidal model in which those timescales can differ by orders of magnitude. In his scenario the planets would first align the stellar rotation, and only much later speed up the rotation and spiral inward. In this respect it is interesting that the tidal timescale calibrated via synchronization apparently sorts the systems consistently relative to each other, but the overall timescale is too long by orders of magnitude. As mentioned above we divided the timescales displayed in Figure~\ref{fig:tau} by $5\cdot 10^{9}$. This could imply that realignment happens on a shorter timescale than synchronization. However the calibration of the synchronization timescale was done with binary star systems having $q\approx 1$, and the tidal mechanism itself might be different for different regimes of its strength \citep{weinberg2012}. One might be able to test the hypothesis of \cite{lai2012} by seeking evidence for excess rotation in stars that are thought to have been realigned ($\lambda \approx 0^\circ$). This could be done by measuring the stellar rotation period or $v \sin i_{\star}$ (if one is willing to assume $\sin i_\star$ is near unity in such systems). If an age estimate is also available, then one could employ the same approach as \cite{schlaufman2010} to assess whether or not the star is rotating at a typical rate, or if it is in the process of being spun up by the planet. This type of analysis was pursued by \cite{pont2009}, though not with this specific hypothesis test in mind. This analysis could be profitably revisited now that many more systems are available for study. There is one caveat to this approach, which is that stars undergo very rapid spin evolution early in their lives due to disk-locking and magnetic braking, i.e., for reasons unrelated to planets. If hot Jupiters arrive very early in the star's history, the realignment might happen in an epoch of rapid decrease of stellar rotation and any memory of an increased rotation due to tides raised by the planet might be lost. The recent work by \cite{hansen2012} presented a calibration of the equilibrium tide theory using the measured parameters of hot-Jupiter systems. While the approach taken in his work is different from ours, he arrived at similar conclusions to those described here: tides have shaped the obliquity distribution in these systems, and there is no theoretical need for core-envelope decoupling. \subsection{High obliquities: a result of dynamical interactions, or initially inclined disks?} \label{sec:cause} To interpret the finding that the host stars of hot Jupiters once had a broad distribution of obliquities, we must answer a crucial question. We need to know if the original obliquity is related to the existence of the hot Jupiter, or if stars and their protoplanetary disks are frequently misaligned for reasons unrelated to hot Jupiters. One might expect an initially close alignment between a star and its protoplanetary disk, as is observed in the Solar system and has been generally assumed in the exoplanet literature. However this is not a foregone conclusion, and indeed several authors have recently challenged this assumption, proposing that the Sun's low obliquity may be an atypical case. \cite{bate2010} proposed that a disk might become inclined with respect to the rotation axis of the central star, as a result of the complex accretion environment within a star cluster. In such a dense environment the tidal interaction with a companion star or other nearby stars could produce chaotic perturbations in the orbits of infalling material, with the material accreting later (destined to become planets) having a different orientation than the material that accreted earlier onto the star. \cite{thies2011} studied the process of inclined infall of gas in detail and found that short period planets on eccentric and inclined orbits can be created in this way. A completely different mechanism for generating primordial misalignments was proposed by \cite{lai2011b}, relying on a magnetic interaction between a young star and the inner edge of its accretion disk. In these scenarios, the star has a high obliquity even though the planets may have never left the plane of the disk out of which they have formed, and therefore the measurements of obliquities bear information about the processes surrounding star formation rather than planet migration. How can one distinguish between misalignment created during the time the planet is still embedded in the disk or after the disk dissipated? One approach, pursued by \cite{watson2011}, is to assess the degree of alignment between stars and their debris disks. Assuming that the stars as well as the debris disks trace the alignment of their predecessors one would learn about the degree of the alignment during the final stages of planet formation. \cite{watson2011} found the inclinations of debris disks and their stars in a sample of 9 systems to be consistent with good alignment. They caution that in their study only systems with $T_{\rm eff}< 6140$~K have been observed and that misaligned system are found around hotter stars. However as we have argued above, the found low obliquities that prevail around cool stars may be a consequence of tidal evolution and not of the mechanism which creates the obliquities. Another approach is to measure the obliquities in binary star systems. If disks would be tilted relative to stellar rotation due to close encounters, then this could also lead to tilted rotation axes in double star systems. There should also be a trend towards misalignment with larger separation between the components in these systems. Conducting such measurements and seeking evidence for such trends is one of the goals of an ongoing observational program entitled BANANA \citep[Binaries Are Not Always Neatly Aligned;][]{albrecht2011}. A more direct way to answer the question raised in the preceding section has recently become possible, thanks to the discovery of systems with multiple transiting planets. A number of arguments support the idea that the orbital planes in such systems are closely aligned; most recently \cite{fabrycky2012} used the measured transit durations to show that the typical mutual inclinations are of order $2^{\circ}$. Therefore it is reasonable to assume that the multiple planetary orbits trace the plane in which the planets formed. Any disruptive dynamical interactions, such as those which have been proposed to explain hot Jupiters, would likely have produced higher mutual inclinations in the multiple-transiting systems. Under that assumption, RM measurements (or other measures of obliquity) in those multiple-transiting systems would establish the angle between the circumstellar disk and the stellar equator. If good alignment is found to be the rule, then the high obliquities in hot Jupiter systems would be more readily interpreted as a consequence of planet migration than as primordial star-disk misalignments. In this paper we have presented new observations of the RM effect for 14 systems harboring hot Jupiters. In addition we critically reviewed the literature, in some cases re-analyzing data that had been obtained previously in order to conform with our protocols. We then used these data to show that the distribution in obliquities is consistent with being shaped by tides raised by the hot Jupiters on the stars. For this we revisited the correlation between the projected obliquity and the effective temperature discovered by \cite{winn2010}, now with a sample of RM measurements twice as large as was then available. We showed that the new measurements agree with the pattern proposed by \cite{winn2010}. With the enlarged sample we showed that obliquity in systems with close in massive planets further depend on the mass ratio and the distance between star and planet, in roughly the manner expected if tides are responsible for the low obliquities. Motivated by these results we then devised two different parameters that represent, at least crudely, the theoretical tidal timescales. This showed that systems which are expected to align fast are all showing projections of the obliquities which are consistent with good alignment. In contrast, systems for which tidal interaction is expected to be weak, due to the stellar structure, distance, or mass ratio, show a nearly random distribution of projected obliquities. Our interpretation is that stars with hot Jupiters once had a very broad range of obliquities. It is tempting to argue further that the large obliquities originate from the same process that produces hot Jupiters, thereby favoring explanations involving dynamical scattering or the Kozai effect, and disfavoring the gradual inspiral due to torques in a protoplanetary disk. However, more observations are needed to check on the possibility that stars and their disks are frequently misaligned for reasons unrelated to hot Jupiters. Among these observations are the extension of RM studies to planets other than hot Jupiters, and measurements of obliquities in binary star systems and in systems with multiple transiting planets. | 12 | 6 | 1206.6105 |
1206 | 1206.0179_arXiv.txt | We investigate a spatially flat Friedmann-Robertson-Walker (FRW) cosmological model with cold dark matter coupled to a modified holographic Ricci dark energy through a general interaction term linear in the energy densities of dark matter and dark energy, the total energy density and its derivative \cite{Chimento:2011dw}. Using the statistical method of $\chi^2$-function for the Hubble data, we obtain $H_0=73.6$km/sMpc, $\omega_s=-0.842$ for the asymptotic equation of state and $ z_{acc}= 0.89 $. The estimated values of $\Omega_{c0}$ which fulfill the current observational bounds corresponds to a dark energy density varying in the range $0.25R < \ro_x < 0.27R$. | Many different observational sources such as the Supernovae Ia \cite{astro-ph/9805201}-\cite{astro-ph/9812133}, the large scale structure from the Sloan Digital Sky survey \cite{arXiv:0707.3413} and the cosmic microwave background anisotropies \cite{arXiv:1001.4538} have corroborated that our universe is currently undergoing an accelerated phase. The cause of this behavior has been attributed to a mysterious component called dark energy and several candidates have been proposed to fulfill this role. For example, a positive cosmological constant $\Lambda$, explains very well the accelerated behavior but it has a deep mismatch with the theoretical value predicted by the quantum field theory. Another issue of debate refers to the coincidence problem, namely: why the dark energy and dark matter energy densities happen to be of the same order precisely today. In order to overcome both problems, it has proposed a dynamical framework in which the dark energy varies with the cosmic time. This proposal has led to a great variety of dark energy models such as quintessence \cite{astro-ph/9807002}, exotic quintessence \cite{arXiv:0706.4142}, N--quintom \cite{arXiv:0811.3643} and the holographic dark energy (HDE) models \cite{hep-th/0403127} based in an application of the holographic principle to the cosmology. According to this principle, the entropy of a system does not scale with its volume but with its surface area and so in cosmological context will set an upper bound on the entropy of the universe \cite{hep-th/9806039}. It has been suggested \cite{hep-th/9803132} that in quantum field theory a short distance cut-off is related to a long distance cut-off (infra-red cut-off L) due to the limit set by the formation of a black hole. Further, if the quantum zero-point energy density caused by a short distance cut-off is taken as the dark energy density in a region of size L, it should not exceed black hole mass of the same size, so $\rho_{\Lambda}=3c^{2}M^{2}_{~P}L^{-2}$, where $c$ is a numerical factor. In the cosmological context, the size L is usually taken as the large scale of the universe, thus Hubble horizon, particle horizon, event horizon or generalized IR cutoff. Among all the interesting holographic dark energy models proposed so far, here we focus our attention on a modified version of the well known Ricci scalar cutoff \cite{arXiv:0810.3663}. Besides, there could be a hidden non-gravitational coupling between the dark matter and dark energy without violating current observational constraints and thus it is interesting to develop ways of testing an interaction in the dark sector. Interaction within the dark sector has been studied mainly as a mechanism to solve the coincidence problem. We will consider an exchange of energy or interaction between dark matter and dark energy which is a linear combination of the dark energy density $\ro_x$, total energy density $\ro$, dark matter energy density $\ro_c$, and the first derivate of the total energy density $\ro'$. \cite{arXiv:0911.5687} | We have examined a modified holographic Ricci dark energy coupled with cold dark matter and found that this scenario describes satisfactorily the behavior of the energy densities of both dark components alleviating the problem of the cosmic coincidence. We have shown that the compatibility between the modified and the global conservation equations restricts the equation of state of the dark energy component relating it to the ratio of energy densities. This constrain makes the holographic density always interacts with the non-holographic component except in the unlikely event that $\alpha = 1$, which is forbidden for positive energy densities. From the observational point of view we have obtained the best fit values of the cosmological parameters $z_{acc}=0.89$, $H_0=73.6$km/sMpc \ and $\gamma_s=0.158$ with a $\chi^{2}_{dof}=0.761 < 1$ per degree of freedom. The $H_{0}$ value is in agreement with the reported in the literature \cite{Riess:2009pu} and the critical redshift $z_{acc}=0.89$ is consistent with BAO and CMB data \cite{Li:2010da}. We have found that in the redshift interval where is trustworthy compared with old stellar sources the model is free from the cosmic-age problem. | 12 | 6 | 1206.0179 |
1206 | 1206.0665_arXiv.txt | We present a measurement of the volumetric Type Ia supernova (SN~Ia) rate (\snr) as a function of redshift for the first four years of data from the Canada-France-Hawaii Telescope (CFHT) Supernova Legacy Survey (SNLS). This analysis includes $286$ spectroscopically confirmed and more than $400$ additional photometrically identified SNe Ia within the redshift range $0.1\leq z\leq 1.1$. The volumetric \snr\ evolution is consistent with a rise to $z\sim 1.0$ that follows a power-law of the form (1+$z$)$^{\alpha}$, with $\alpha={2.11\pm 0.28}$. This evolutionary trend in the SNLS rates is slightly shallower than that of the cosmic star-formation history over the same redshift range. We combine the SNLS rate measurements with those from other surveys that complement the SNLS redshift range, and fit various simple SN Ia delay-time distribution (DTD) models to the combined data. A simple power-law model for the DTD (i.e., $\propto t^{-\beta}$) yields values from $\beta=0.98\pm0.05$ to $\beta=1.15\pm0.08$ depending on the parameterization of the cosmic star formation history. A two-component model, where \snr\ is dependent on stellar mass (\mstellar) and star formation rate (SFR) as $\snr(z)=A\times \mstellar(z) + B\times\mathrm{SFR}(z)$, yields the coefficients $A=(1.9\pm 0.1)\Aunit$ and $B=(3.3\pm 0.2)\Bunit$. More general two-component models also fit the data well, but single Gaussian or exponential DTDs provide significantly poorer matches. Finally, we split the SNLS sample into two populations by the light curve width (stretch), and show that the general behavior in the rates of faster-declining SNe~Ia ($0.8\leq s < 1.0$) is similar, within our measurement errors, to that of the slower objects ($1.0\leq s < 1.3$) out to $z\sim 0.8$. | Type Ia supernova (SN~Ia) explosions play a critical role in regulating chemical evolution through the cycling of matter in galaxies. As supernovae (SNe) are the primary contributors of heavy elements in the universe, observed variations in their rates with redshift provide a diagnostic of metal enrichment over a cosmological timeline. The frequency of these events and the processes involved provide important constraints on theories of stellar evolution. SNe~Ia are thought to originate from the thermonuclear explosion of carbon-oxygen white dwarfs that approach the Chandrasekhar mass via accretion of material from a binary companion \citep[for reviews, see][]{hn00,how11}. This process can result in a significant ``delay time'' between star formation and SN explosion, depending on the nature of the progenitor system \citep{mad98,gre05}. The SN~Ia volumetric rate (\snr) evolution therefore represents a convolution of the cosmic star-formation history with a delay-time distribution (DTD). As such, measuring the global rate of SN~Ia events as a function of redshift may be useful for constraining possible DTDs and, ultimately, progenitor models -- the detailed physics of SNe Ia remains poorly understood, with several possible evolutionary paths \citep[e.g.][]{bra95,Liv00}. One complication for rates studies is that many SN surveys at low redshifts are galaxy-targeted, counting discoveries in a select sample of galaxies and converting to a volumetric rate by assuming a galaxy luminosity function. This method can be susceptible to systematic errors if it preferentially samples the bright end of the galaxy luminosity function, biasing toward SNe in more massive, or brighter, galaxies \citep[see, e.g.,][]{sul10}. Since many SN Ia properties are correlated with their hosts, the recovered rates may then not be representative of all types of SNe~Ia. A second type of SN survey involves making repeat observations of pre-defined fields in a ``rolling search'', to find and follow SNe in specific volumes of sky over a period of time. Such surveys minimize the influence of host bias, but still suffer from Malmquist bias and other selection effects. It is reasonably straight forward --- although often computationally expensive --- to compensate for the observational biases within rolling searches. The advent of these wide-field rolling surveys has significantly enhanced SN statistics at cosmological distances. The Supernova Legacy Survey (SNLS) in particular has contributed a large sample of Type Ia SNe out to redshifts of $z\sim1.05$ \citep{guy10}. Although its primary goal is to assemble a sample of SNe~Ia to constrain the cosmological parameters \citep[e.g.][]{ast06,sul11}, the SNLS is also ideal for studies of SN rates \citep[][]{nei06,baz09}. The SNLS is a rolling high-redshift search, with repeat multi-color imaging in four target fields over five years and as such has consistent and well-defined survey characteristics, along with significant follow-up spectroscopy. However, due to the selection effects (including incomplete spectroscopic follow-up) and other systematic errors, such as contamination and photometric redshift errors, present in any SN survey, a detailed understanding of internal biases is necessary for accurate rate calculations. \begin{figure} \plotone{f01.eps} \caption{Volumetric SN~Ia rates as a function of redshift from various previous studies, taken from \citet{li11b,dil10,rod10,dah08,gra11,nei06}. Additional individual rates ($+$) include, in order of increasing redshift: \citet{bla04,bot08,kuz08}. Values are plotted as published, with the exception of a correction to the cosmology used in this paper. As a comparison, the lines shows the evolution of various model cosmic star-formation histories from \citet[][piece-wise fit is the short-dashed line, the \citet{col01} form is the long-dashed line]{li08} and \citet[][dot-dashed line]{yuk08}.} \label{fig:rates_lit} \end{figure} In the past decade, volumetric SN~Ia rates have been measured to varying degrees of accuracy out to redshifts of $z\sim 1.6$ (Fig.~\ref{fig:rates_lit}). \citet{cap99} compute the SN~Ia rate in the local universe ($z\sim 0.01$) from a combined visual and photographic sample of $\sim 10^4$ galaxies, yet their ability to distinguish core-collapse SNe from Type Ia SNe was severely limited. More recent work by \citet{li11b} using $\sim270$ SNe Ia from the Lick Observatory Supernova Search \citep[LOSS;][]{lea11} has made significant improvements in the statistics over previous studies on local SNe~Ia. The rates published by \citet{dil10} include data from 516 SNe~Ia at redshifts $z<0.3$ from the SDSS-II Supernova Survey (SDSS-SN), with roughly half of these confirmed through spectroscopy. At intermediate redshifts, rate measurements are provided by \citet[][38 SNe from the Supernova Cosmology Project in the range $0.25\leq z \leq 0.85$]{pai02}, \citet[][8 SNe within $0.3<z<1.2$]{ton03}, and \citet[][$>100$ SNe from the IfA Deep Survey, 23 of which have spectra]{rod10}. \citet{nei06} used a spectroscopic sample of 58 SNe Ia from the first two years of SNLS to measure a cumulative volumetric rate in the redshift range $0.2<z<0.6$. SN~Ia rates out to $z\sim1.6$ are presented by \citet{dah04} using 25 SNe Ia (19 with spectra) from \textit{Hubble Space Telescope (HST)} observations of the Great Observatories Origins Deep Survey (GOODS) fields. These data were reanalyzed by \citet{kuz08} using a Bayesian identification algorithm, and the \textit{HST} sample updated by \citet{dah08} extending the 2004 sample to 56 SNe. Ground-based measurements from the Subaru Deep Field have also been made by \citet{poz07} using 22 SNe~Ia, updated by \citet{gra11} with 150 events. The general trend of Fig.~\ref{fig:rates_lit} reveals that the rates typically increase from $z=0$ to $z=1$. There is a wide spread in the existing rate measurements, particularly in the range $0.4 < z < 0.8$. At higher redshifts, data from the GOODS collaboration provide some apparent evidence for a turnover in the SN~Ia rates. In particular, \citet{dah04,dah08} report a decline in SN~Ia rates beyond $z\sim 0.8$. If present, this decline might point to a larger characteristic delay time between star formation and SN explosion \citep[see also][]{str04}. However, another independent analysis of the \textit{HST} GOODS data finds rates that are offset, with measurements by \citet{kuz08} consistently lower than those of \citet{dah04,dah08}. \citet{kuz08} argue that their results do not distinguish between a flat or peaked rate evolution. Ground-based data in this range \citep{gra11}, while consistent with the \textit{HST}-based results, show no obvious evidence for a decline above $z\sim1$. In this paper we use four years of data from the SNLS sample to investigate the evolution of SN~Ia rates with redshift out to $z\sim 1.1$. The sample presented comprises $\sim 700$ photometrically identified SNe~Ia from SNLS detected with the real-time analysis pipeline \citep{per10}. One third of these have been typed spectroscopically, and one half of the $\sim700$ have a spectroscopic redshift (sometimes from ancillary redshift surveys in the SNLS fields). No other data set currently provides such a well-observed and homogeneous sample over this range in redshift. Additionally, rigorous computation of the survey detection efficiencies and enhancements in photometric classification techniques are incorporated into the new SNLS rate measurements. Monte Carlo simulations of artificial SNe~Ia with a range of intrinsic parameters are performed on all of the detection images used in the SNLS real-time discovery \citep{per10}; these provide an exhaustive collection of recovery statistics, thereby helping to minimize the effects of systematic errors in the rate measurements. The SNLS SNe~Ia can be used to examine the relationship between the \snr\ and redshift, given some model of the SN Ia DTD. The size of the SNLS sample also permits a division of the SNe Ia by light-curve width \citep[in particular the ``stretch''; see][]{per97}, allowing a search for differences in the volumetric rate evolution expected by any changing demographic in the SN Ia population. Brighter, more slowly-declining (i.e., higher stretch) SNe~Ia are more frequently found in star-forming spirals, whereas fainter, faster-declining SNe~Ia tend to occur in older stellar populations with little or no star formation \citep{ham95,sul06b}. If the delay time for the formation of the lowest-stretch SNe~Ia is sufficiently long (i.e., their progenitors are low-mass stars $\sim 10$ Gyr old), these SNe~Ia will not occur at high redshifts \citep{how01}. The behavior of the high-$z$ rates can reveal the properties of the progenitor systems. The organization of this paper is as follows: An overview of the rate calculation is provided in \S\ref{sec:ratecalc}. The SNLS data set, along with the light-curve fitting and selection cuts used to define the photometric sample, is introduced in \S\ref{sec:SNLS}. SN~Ia detection efficiencies and the rate measurements are presented in \S\ref{sec:effs} and \S\ref{sec:SNLSrates}, respectively. Several models of the SN Ia DTD are then fit to the rate evolution in \S\ref{sec:dtds}, and the results discussed. Finally, the stretch dependence of the rate evolution is investigated in \S\ref{sec:stretchdep}. We adopt a flat cosmology with ($\Omega_M$,$\Omega_\Lambda$)=(0.27,0.73) and a Hubble constant of $H_0=70\,\mathrm{km}\,\mathrm{s}^{-1}\,\mathrm{Mpc}^{-1}$. | In this paper, we have probed the volumetric rate evolution of ``normal'' $0.8<s<1.3$ SNe Ia using a sample of $691$ events from the Supernova Legacy Survey (SNLS) in the range $0.1<z<1.1$, $286$ of which have been confirmed spectroscopically. The SNLS rates increase with redshift as (1+$z$)$^{\alpha}$ with $\alpha={2.11\pm 0.28}$, and show no evidence of flattening beyond $z\sim 0.5$. Due to spectroscopic incompleteness and the decrease in detection efficiency for the SNLS sample, a rollover in the slope cannot be ruled out beyond $z\sim 1$ based on the SNLS data alone. As a significant component of the SN~Ia rate is linked with young stellar populations, an increasing fraction of SN~Ia events may suffer the effects of host extinction at higher redshifts. In our rate calculation method, the effect of SN color is factored directly into the detection efficiency determinations: detection recovery is evaluated empirically according to the observed SN color regardless of its cause. Redder objects at a given redshift have lower detection efficiencies, and are correspondingly more heavily weighted in the rates determination. Combining the SNLS data with that from other SN Ia surveys, we fit various simple delay-time distributions (DTDs) to the volumetric SN Ia rate data. DTDs with a single Gaussian are not favored by the data. We find that simple power-law DTDs ($\Psi(t)\propto t^{-\beta}$) with $\beta\sim1$ ($\beta=0.98\pm0.05$ to $\beta=1.15\pm0.08$ depending on the parameterization of the cosmic SFH) can adequately explain all the SN Ia volumetric rate data, as can two-component models with a prompt and delayed channel. These models cannot be separated with the current volumetric rate data. Integrating these different DTDs gives the total number of SNe Ia per solar mass formed (excluding sub-luminous $s<0.8$ events) of $N_{\mathrm{Ia}}/\mstar\sim4.4-5.7\times10^{-4}\,\mathrm{SNe}\,\msun^{-1}$ (assuming a Salpeter IMF), depending on the star formation history and DTD model. This is in good agreement with other similar analyses, but lies significantly below the number expected from DTDs derived from cluster SN Ia rates. The use of other techniques, such as fitting the SFH of individual galaxies \citep{sul06b,bra10,mao11}, or observing a simplified subset of galaxies \citep{tot08,mao10}, use more information, and in principle ought to be more reliable. However, each technique has significant drawbacks, such as contamination \citep{tot08,mao11}, limitations of SED fitting codes \citep{sul06b,bra10,mao11}, and the assumption that all cluster galaxies formed at $z=3$ in a delta-function of star-formation \citep{mao11}. Therefore, our results are an important complementary constraint. By presenting an evolution in the SN Ia rate over a large redshift baseline done self-consistently by a single survey we have for the first time mitigated the primary drawback of this method -- having to combine myriad rate determinations from multiple surveys, all done with different assumptions and biases, sometimes disparate by large factors \citep{nei06}. We also find no clear evidence for a difference in the rate evolution for SNLS samples with $0.8\leq s < 1.0$ and $1.0\leq s< 1.3$ out to $z=0.8$, although the stretch evolution model from \citet{how07} cannot be ruled out conclusively. Stretch evolution plays a more significant role in the sub-luminous population \citep{gon11}, which show a much flatter evolution than the $s>0.8$ sample. Next generation surveys such as Dark Energy Survey (DES), Pan-STARRS, Palomar Transient Factory (PTF), and SkyMapper, many of which are already underway, are finding thousands of SNe Ia (in comparison to the $\sim 700$ in this study). Statistical rate determinations ought to improve, but systematic difficulties will remain, as not all SNe can be spectroscopically confirmed. However, large number statistics will allow the construction of sub-samples larger than the three (split by stretch) analyzed here. Comparison of the relative rates of SNe with different properties and in different environments may ultimately improve deduced DTDs, and allow for the construction of different DTDs for subsets of SNe Ia. | 12 | 6 | 1206.0665 |
1206 | 1206.7012.txt | We present the results of the first spectroscopic follow-up of 132 optically blue UV-excess sources selected from the UV-excess survey of the Northern Galactic Plane (\UVEX). The UV-excess spectra are classified into different populations and grids of model spectra are fit to determine spectral types, temperatures, surface gravities and reddening. From this initial spectroscopic follow-up 95$\%$ of the UV-excess candidates turn out to be genuine UV-excess sources such as white dwarfs, white dwarf binaries, subdwarfs type O and B, emission line stars and QSOs. The remaining sources are classified as slightly reddened main-sequence stars with spectral types later than A0V. The fraction of DA white dwarfs is 47$\%$ with reddening smaller than $E(B-V)$$\leq$0.7 mag. Relations between the different populations and their \UVEX photometry, Galactic latitude and reddening are shown. A larger fraction of \UVEX white dwarfs is found at magnitudes fainter than $g$$>$17 and Galactic latitude smaller than $|b|<$4 compared to main-sequence stars, blue horizontal branch stars and subdwarfs. | Traditionally, surveys searching for faint blue objects have avoided the Galactic Plane because of the high dust absorption. Surveys searching for quasars and white dwarfs therefore mostly observed at Galactic latitudes larger than $|b|>$30\degr. Examples of such surveys are the Palomar Green survey (PG, Green et al., 1986), the Kiso survey (Wegner et al., 1987, Limoges et al., 2010), the Sloan Digital Sky survey (SDSS, York et al., 2000, Yanni et al., 2009 and Eisenstein et al., 2006) and the Hamburg Quasar survey (HQS, Hagen et al., 1995, Homeier et al., 1998) in the northern hemisphere and the Montreal-Cambridge-Tololo survey (MCT, Lamontagne et al., 2000, Demers 1986), the Edinburgh-Cape survey (EC, Kilkenny et al., 1997, Stobie et al., 1997), the Homogeneous Bright Quasar survey (Gemmo et al., 1995) and the Hamburg-ESO survey (Christlieb et al., 2001, Wisotzki et al., 1996) in the southern hemisphere. Only the Kitt Peak-Downes survey (KPD, Downes et al., 1986) survey and the Sandage Two-colour Galactic Plane survey (Lanning, 1973) observed a bit closer to the Galactic Plane. Some of the brighest \UVEX UV-excess sources are Lanning sources (e.g. UVEXJ0328+5035 and UVEXJ0528+2716 in Table\ \ref{tab:spectra} are in Lanning, 1973 and Lanning et al., 2004 respectively). The lowest Galactic latitudes $|b|<5\degr$ are still relatively unexplored (see e.g. Fig.2 of Napiwotzki et al., 2003). In order to determine key population characteristics of Galactic sources, such as their scaleheight or space density, it is crucial to study the low Galactic latitude environment. The space density of stellar remnants, such as white dwarfs, Cataclysmic Variables and AM CVn stars, is currently poorly constrained while there must be $\sim10^{5}$ of them in the Milky Way (see Fig.\,1 of Groot et al., 2009, McCook et al., 1999, L\'epine et al., 2011 and Nelemans et al., 2001).\\ One of the main goals of the European Galactic Plane Surveys (\EGAPS) is to obtain a homogeneous sample of evolved objects in our Milky Way with well-known selection limits. The UV-excess survey of the Northern Galactic Plane (\UVEX, Groot et al., 2009) images a 10$\times$185 degrees wide band (--5\degr$<$ $b$ $<$+5\degr) centred on the Galactic equator in the $U,g,r$ and $\he$ bands down to $\sim 21^{st}-22^{nd}$ magnitude using the Wide Field Camera mounted on the Isaac Newton Telescope on La Palma. From the first 211 square degrees of \UVEX data, a catalogue of 2\,170 optically blue UV-excess candidates was selected in Verbeek et al. (2012; hereafter V12). These UV-excess sources were selected from the $(U-g)$ versus $(g-r)$ colour-colour diagram and $g$ versus $(U-g)$ and $g$ versus $(g-r)$ colour-magnitude diagrams by an automated field-to-field selection algorithm. This automated selection algorithm and the properties of the selected UV-excess catalogue are described in V12. Less than $\sim$1$\%$ of the selected UV-excess sources are currently known in the literature.\\ Here we report our spectroscopic follow-up for 132 objects (6$\%$) in the UV-excess catalogue of V12. This early reconnaissance is important for the design of future colour-selection methods for various populations, comparable to the selection techniques for e.g. the SDSS, which generally do not have to deal with the added complication of reddening (G\"ansicke et al., 2009, Girven et al., 2011). In Sect.\ \ref{sec:followup} the spectroscopy of the selected sample is described, and in Sect.\ \ref{sec:spectra} the spectra are presented and classified. The spectra are fitted to grids of model spectra in order to determine the characteristics of UV-excess spectra classified as white dwarfs, subdwarfs, main-sequence stars and blue horizontal branch stars. Finally in Sect.\ \ref{sec:discussion} we summarise the conclusions of the UV-excess catalogue and the spectroscopic follow-up. The UV-excess spectra are shown in Figs.\ \ref{fig:spectra1} to\ \ref{fig:spectra11} and their features are listed in Table\ \ref{tab:spectra} in Appendix\ \ref{app:appendix}. All spectra and the table can also be obtained from the \UVEX website\footnote{http://www.uvexsurvey.org}.\\ \begin{figure*} \centerline{\epsfig{file=Classification2CCDs25mei2012.eps,width=30cm,angle=0,clip=}} \caption{The UVEX colour-colour diagrams with the classified UV-excess candidates. The lines are the simulated colours of unreddened main-sequence stars (solid black) and the O5V-reddening line (dashed black) of V12. The cyan and green dashed lines are respectively the simulated colours of unreddened Koester DA and DB white dwarfs. The different symbols indicate the classification: White Dwarf (DA/DB/DAB/DC/DZ/DAe), White Dwarf+Red Dwarf binary (DA+dM), Cataclysmic Variable (CV), T Tauri star (TT), Be star (Be), subdwarf star (sdO/sdB), main-sequence star or blue horizontal branch star (MS/BHB), G2V star and M-giant (G/M), Quasi Stellar Object (QSO) and unknown (?). The sources classified as ``noisy'' in Sect.\ \ref{sec:mainseqsubdwarf} are not shown here. There is one more H$\alpha$ emitter classified as T Tauri star at $(g-r)$=1.55, $(HeI-r)$=0.6 and one M-giant at $(g-r)$=0.28, $(HeI-r)$=--1.9 not shown in the $(HeI-r)$ vs. $(g-r)$ colour-colour diagram. \label{fig:classCCD}} \end{figure*} \begin{figure*} \centerline{\epsfig{file=Classification2CMDs25mei2012.eps,width=24cm,angle=0,clip=}} \caption{The UVEX colour-magnitude diagrams with the classified UV-excess candidates. \label{fig:classCMD}} \end{figure*} \begin{figure*} \centerline{\epsfig{file=ClassificationIPHASCCDandCMD25mei.eps,width=24cm,angle=0,clip=}} \caption{IPHAS colour-colour and colour-magnitude diagrams with the classified UV-excess candidates that have a match in IPHAS. There is one extra sources classified as M-giant at $(r-i)$=0.1, $(r-H\alpha)$=2.25 in the $(r-H\alpha)$ vs. $(r-i)$ colour-colour diagram. The lines are the synthetic colours of main-sequence stars (black) with reddening $E(B-V)$=0 and $E(B-V)$=1 and unreddened Koester DA white dwarfs (cyan). Sources that are in the Deacon IPHAS-POSSI PM catalogue are encircled blue, sources that are in the Witham H$\alpha$ emission line catalogue are encircled red and sources that show H$\alpha$ emission lines in their spectra but are not in the Witham catalogue are encircled green. The Witham catalogue covers the magnitude range 13$<r<$19.5 and the Deacon catalogue covers the magnitude range 13.5$<r<$19. The sources classified as ``noisy'' in Sect.\ \ref{sec:mainseqsubdwarf} are not shown here, one of them has a match in the IPHAS-POSSI PM catalogue. \label{fig:IPHASCCD}} \end{figure*} | \label{sec:discussion} The main conclusion is that of the sources in the UV-excess catalogue 95$\%$ are genuine UV-excess sources, such as white dwarfs, white dwarf binaries, subdwarf stars type O and B and QSOs. Five percent of the UV-excess candidates are classified as main-sequence (MS) stars or blue horizontal branch (BHB) stars with spectral types later than A0V. If the sources classified as MS/BHB are main-sequence stars, the fitting of the Pickles library spectra (Pickles, 1998) shows that 4 sources are slightly reddened F0V stars with reddening $E(B-V)$$\leq$0.1 and 3 sources are B0V-B3V stars with reddening 0.4$\leq$$E(B-V)$$\leq$0.5. Their spectra look like B-type main-sequence stars, but the Balmer absorption lines are stronger than the Balmer lines of the best fit template spectra. Since gravity is slightly high they could be horizontal branch stars, although usually BHB stars have less helium absorption. Two of the B-type MS/BHB stars slightly blue shifted so might be high velocity stars, UVEXJ1940+3220 has a velocity of $RV$=320km/s and UVEXJ1938+3054 has a velocity of $RV$=163km/s. There is 1 G-type star with as best fit a Pickles G2V star with reddening $E(B-V)$$\sim$0.3. G-type stars are expected to have colours redder than $(g-r)$$>$0.6 and $(U-g)$$>$0.4, but they can enter the UV-excess region when they are metal weak, i.e. subdwarfs type with less light blocked at the blue/UV wavelengths (Eracleous et al., 2002). Since the number of late type stars in the fields is large, photometric errors can cause a few outliers to be scattered into the UV-excess selection region (Krzesinski et al., 2004).\\ Secondly, in the colour-colour and colour-magnitude diagrams of Figs.\,6 and 7 of V12 about 20$\%$ of the UV-excess sources overlaps with the location of the `subdwarfs' population at $(g-r)$$>$0.3 and $(U-g)$$>$0.2. These UV-excess sources that overlap the subdwarf area in the colour-colour and colour-magnitude diagrams are 2 QSOs, 7 Cataclysmic Variables, 1 DAe, 2 T Tauri stars, 1 Be star, 3 DA+dM stars, 3 He-sdO stars, 1 sdO star, 2 sdB stars, 1 B-type MS/BHB star, 3 F-type MS/BHB stars and 1 DAB white dwarf. For our selection of UV-excess candidates from the \UVEX data this means that when the aim is to find white dwarfs, a colour cut can be applied to decrease the number of other objects. But leaving these sources at the location of the `subdwarfs' out from the UV-excess catalogue would lead to a loss of most QSOs, Cataclysmic Variables, T Tauri stars, Be stars and DA+dM stars.\\ The location of the different populations in the $(U-g)$ vs. $(g-r)$ colour-colour diagram and the $g$ vs. $(U-g)$ and $g$ vs. $(g-r)$ colour-magnitude diagrams is shown in Fig.\ \ref{fig:classCCD} and\ \ref{fig:classCMD}. Their positions match with the positions of the populations in the colour-colour diagrams of other surveys (e.g. Fig.\,1 of Krzesinski et al., 2004, Fig.\,1 of Harris et al., 2003, Fig.\,3 of Stobie et al., 1997, Fig.\,3 of Yanny et al., 2009 and the Figs. of Kilkenny et al., 1997). The locations of the classified sources in the colour-colour and colour-magnitude diagrams agree with the locations of the sources with a Simbad match in the colour-colour and colour-magnitude diagrams in Fig.\,9 of V12. There is a clear relation between the different kind of sources and the way they are selected in V12. The way the sources were selected from the colour-colour and colour-magnitude diagrams is captured in the `selection label' (column 20 of the UV-excess catalogue, Appendix A of V12), and is summarized for our classified sources in Table.\ \ref{tab:selection}. Only 2 DA white dwarfs were selected less than 0.4 magnitude from the blue edge in the $g$ vs. $(g-r)$ colour-magnitude diagram, the other 60 DA white dwarfs were selected more than 0.4 magnitude from the blue edge in the $g$ vs. $(g-r)$ colour-magnitude diagram. All DB and DC white dwarfs were selected both were selected more than 0.4 magnitude from the blue edge in both colour-magnitude diagrams and in the $(U-g)$ vs. $(g-r)$ colour-colour diagram, while most DBA white dwarfs were selected less than 0.4 magnitude from the blue edge in the $g$ vs. $(g-r)$ colour-magnitude diagram and in the $(U-g)$ vs. $(g-r)$ colour-colour diagram. The 2 QSOs and the majority of the H$\alpha$ emission line objects were selected in the $g$ vs. $(U-g)$ colour-magnitude diagram but not in the $g$ vs. $(g-r)$ colour-magnitude diagram. We could improve the selection method of V12 using these spectroscopic results by taking only sources in the UV-excess catalogue with favourable selection labels into account. This will increase the number of genuine UV-excess objects to 97$\%$ by e.g. leaving out selection labels `514' and `518' since the largest fraction MS/BHB stars have these selection labels, but this will also lead to a loss of some peculiar objects such as DAe stars, Cataclysmic Variables and Be stars. The \UVEX and \IPHAS photometry can also be combined with other (infrared and ultraviolet) surveys in order to improve the selection of different populations (Verbeek et al., in prep.).\\ \begin{table*} \caption[]{The selection in V12 of the classified UV-excess spectra. \label{tab:selection} } {\small \begin{tabular}{ | l | l | l | } \hline Label & Selected from & Objects \\ \hline 514 & $g$ vs. $(U-g)$ & 1DAe, 1Be, 1MS/BHB, 1He-sdO, 1sdB+F \\ 515 & $g$ vs. $(U-g)$ \& $(U-g)$ vs. $(g-r)$ & 2CV, 2QSO, 1TT, 1DAB+dM, 1He-sdO, 1BHB/MS \\ 518 & $g$ vs. $(U-g)$ \& $<$0.4$g$ vs. $(g-r)$ & 1CV, 1DAe, 2MS/BHB, 1noisy \\ 519 & $g$ vs. $(U-g)$ \& $<$0.4$g$ vs. $(g-r)$ \& $(U-g)$ vs. $(g-r)$ & 5CV, 4DBA, 2DA, 1TT, 2DA+dM, 4sdB, 2He-sdO, 1BHB/MS, 1G, 4noisy \\ 1028 & $g$ vs. $(g-r)$ & 28DA, 1DZA, 1unknown, 1DA+dM, 1MS/BHB, 1sdB, 1noisy \\ 1029 & $g$ vs. $(g-r)$ \& $(U-g)$ vs. $(g-r)$ & 3DA, 1DA+dM \\ 1031 & $g$ vs. $(g-r)$ \& $<$0.4$g$ vs. $(U-g)$ \& $(U-g)$ vs. $(g-r)$ & 1DA, 1MIII \\ 1542 & $g$ vs. $(g-r)$ \& $g$ vs. $(U-g)$ & 1unknown \\ 1543 & $g$ vs. $(g-r)$ \& $g$ vs. $(U-g)$ \& $(U-g)$ vs. $(g-r)$ & 28DA, 4DB, 1DBA, 4DC, 1DAe, 1DZ, 1sdO, 1BHB/MS, 6noisy \\ \hline \end{tabular} \\ } \end{table*} About 64$\%$ of the UV-excess candidates turn out to be white dwarfs. The fitting of the white dwarf models to the UV-excess hydrogen atmosphere white dwarf spectra shows a distribution of 9\,000K$<$$T_{\rm eff}$$<$65\,000K and an average surface gravity of $log\,g$$\sim$8. These results are in agreement with the results of other studies. (Liebert et al., 2005, Bergeron et al., 1992, Napiwotzki et al., 1999, Finley et al., 1997, Gianninas et al., 2011 and Kepler et al., 2007). The accuracy of the continuum fitting method applied in Sect.\ \ref{sec:whitedwarf} depends on the SNR and the flux calibration of the spectra. For white dwarfs the accuracy of the temperature fit will approach the surface temperature to typically $\sim$1\,000$K$ for white dwarfs with T$<$20\,000 and $\sim$2\,000$K$ for the hotter white dwarfs with T$\geq$20\,000, for spectra with signal to noise SNR$>$20.\\ When we extrapolate the result that 64$\%$ of the UV-excess catalogue sources are white dwarfs, the complete \UVEX survey will bring up a sample of $\sim 1.2 \times 10^{4}$ new white dwarfs ($\sim$7 per square degree). If we only look at UV-excess white dwarf sample brighter than $g$$<$20, \UVEX will bring up a sample of $\sim$4000 new white dwarfs with $g$$<$20 in the full survey area. % g<20 = 124 sources and g>20 = 8 sources (3DA, DA+dM, DZ, DC, MIII, DAB) --> Total g<20 sample 124 has 78 WDS = 63% % UV-excess sample contains 750 candidates --> 0.63 x 750 x 1850/211 = 4143 The UV-excess sample might not be complete for the coolest white dwarfs below $T$$<$10\,000K since they have too red colours. There is also the additional problem of dust extinction (Sale et al., 2009), which has only a small effect on the local white dwarf sample while it merely screens out more distant objects. As shown in Sect.\ \ref{sec:whitedwarf} reddening is typically $E(B-V)$$\leq$0.1 magnitudes for most of the white dwarfs in the UV-excess catalogue. A space density of white dwarfs (Holberg et al., 2008) in the Galactic Plane and a comparison with population synthesis predictions will be further discussed in Verbeek et al., (in prep.).\\ \begin{figure} \centerline{\epsfig{file=transparent.29jun.eps,width=8cm,angle=0,clip=}} \caption{Galactic latitude distribution of the sources classified as white dwarfs (blue) and the sources classified as sdO/sdB stars, MS/BHB stars and ``noisy'' (red). The number of sources per bin is normalized by the total number of obtained spectra in the latitude bin. \label{fig:latitude}} \end{figure} The Galactic latitude distribution of the sources classified as white dwarfs and as sdO/sdB stars, main-sequence stars and blue horizontal branch stars is shown in Fig.\ \ref{fig:latitude}. The sources labeled as ``noisy'' in Sect.\ \ref{sec:mainseqsubdwarf} are add to the sdO/sdB/BHB/MS sample since they probably are sdO/sdB or MS/BHB stars. The white dwarfs are mainly detected at Galactic latitudes smaller than $|b|$$<$4, while the distribution of sdO/sdB stars and MS/BHB stars peaks at Galactic latitudes larger than $|b|$$>$4. This result can be explained by the absolute magnitude distribution of the different populations in combination with the effect of extinction, as can be seen in Fig.\,1 of Groot et al. (2009).\\ \begin{figure} \centerline{\epsfig{file=transparent.Magnitude29jun.eps,width=8cm,angle=0,clip=}} \caption{Magnitude distribution of the sources classified as white dwarfs (blue) and the sources classified as sdO/sdB stars, MS/BHB stars and ``noisy'' (red). The number of sources per bin is normalized by the total number of obtained spectra in the magnitude bin. \label{fig:magnitude}} \end{figure} The magnitude distribution of the spectra classified as white dwarfs and as sdO/sdB stars, MS/BHB stars and ``noisy'' is shown in Fig.\ \ref{fig:magnitude}. The fraction of white dwarfs clearly increases for fainter $g$-band magnitudes. The total number of white dwarfs increases strongly for fainter magnitudes since also the number of selected sources increases (see e.g. Fig.\,7 of V12 and Fig.\,1 of Bergeron et al., 1992). The fraction of MS/BHB and subdwarf sources is larger for the brighter $g$-band magnitudes, even with the sources classified as ``noisy'' included in the sdO/sdB/BHB/MS sample. A larger fraction of white dwarfs is found at magnitudes fainter than $g$$>$17.\\ \begin{figure*} \centerline{\epsfig{file=GalacticLatitudeVSLongitude25mei.eps,width=18cm,angle=0,clip=}} \caption{Galactic latitude vs. Galactic longitude diagram with all obtained UV-excess spectra. The classified sources are indicated with the symbols of Figs.\,1 to 3. \label{fig:LatitudeVSLongitude}} \end{figure*} If we assume that UV-excess candidates classified as main-sequence stars and blue horizontal branch stars are all MS stars, we can estimate the distance $d$ using $d$=$10^{0.2\times(m-M-A_{V})}\,10pc$, where we use the observed $g$-band magnitude as apparent magnitude ($m$), $M$ is the absolute magnitude and $A(V)$ is the total extinction for the $V$-band filter. Since $A(V)$ = $R_{V}$$\times$$E(B-V)$, where we use $R_{V}$=3.1 for the indicator of dust grain size distribution and the results of the fitting in Sect.\ \ref{sec:mainseqsubdwarf} for the reddening $E(B-V)$=$A(B)-A(V)$, we can estimate a distance range per source. In our UV-excess sample the F-type MS/BHB stars have a $g$-band magnitude of 17.3$<$$g$$<$19.6 and reddening $E(B-V)$=0.1, the G2V star has a $g$-band magnitude of $g$=19.5 and $E(B-V)$=0.3, the B-type BHB/MS stars have a typical $g$-band magnitude between 14.5$<$$g$$<$18.6 and $E(B-V)$=0.4, Taking into account the effect of reddening the distance estimations would be $\sim$5kpc for the G2V star with $g$=19.5 and $\sim$8kpc for the F0V star with $g$=17.3 which is within the Milky Way. %For the faintest F0V star with $g$=19.6 $\sim$20kpc for the. For the fainter F-type and B-type stars the distances would be $\sim$20kpc for the F0V star with $g$=19.6 and $\sim$35kpc for the B0V star with $g$=17.0 if they would be main-sequence stars. These distances would be outside the Milky Way. Since their colours are only slightly reddened they must be intrinsically fainter objects. So, we conclude that these objects must be blue horizontal branch stars or subdwarf type stars.\\ An interesting side benefit is the detection of the 2 broadline QSOs in the UV-excess spectra, with $z$$\sim$2.16 and $z$$\sim$1.48 and at $|b|$=4. Only about $\sim$10 QSOs are found at low Galactic latitude regions (Im et al., 2007, Lee et al., 2008 and Becker et al., 1990). The Schlegel map (Schlegel et al., 1998) gives a reddening of $E(B-V)$$\sim$0.5 for both QSOs. Due to the internal reddening of the QSOs we can not directly estimate the amount of reddening caused by our Milky Way from their spectra (Knigge et al., 2008).\\ Of the UV-excess spectra 122 have a match in \IPHAS. These matches are shown in the colour-colour and colour-magnitude diagrams of Fig.\ \ref{fig:IPHASCCD}. Nine of the classified UV-excess sources are in the Deacon IPHAS-POSSI PM catalogue (Deacon et al., 2009): 7 DA white dwarfs, 1 DC white dwarf and 1 DA+dM binary system. Except for the DA+dM at $(r-i)$=1.3 all sources overlap with the location of the white dwarf population at $(r-i)$$\sim$0 in Fig.\,11 of V12. Eight of the classified UV-excess sources are in the Witham H$\alpha$ emission line catalogue (Witham et al., 2008): 4 Cataclysmic Variables, 1 Be star, 1 Classical T Tauri star and 2 sdO candidates. There are some sources with clear H$\alpha$ emission lines in their spectra that are not in the Witham H$\alpha$ emission line catalogue. Four sources classified as Cataclysmic Variables clearly show H$\alpha$ emission in the \IPHAS colour-colour diagram of Fig.\ \ref{fig:IPHASCCD}. Two of these Cataclysmic Variables are not in the Witham catalogue because they have $r$-band magnitudes $r$$>$19.5. The H$\alpha$ emission of some other Cataclysmic Variables with $EW$$<$20\AA\, is probably not strong enough to be in the Witham catalogue, or they can also have variable emission.\\ \subsection{Comparison with spectroscopic surveys} \label{sec:crossmatch} \begin{itemize} \item We can compare our results with the spectroscopic observations of Eracleous et al. (2002) of 27 UV-bright stars, with $(U-B)$$<$--0.2 and magnitude 13$<$$B$$<$16, selected from the Sandage Two-colour Galactic Plane survey, in the Lanning catalogue (Lanning, 1973). This sample contains 2 DA white dwarfs, 1 DB white dwarfs, 1 DA+dM, 16 O/B stars (60$\%$), 1 F/G star, 1 M star, 1 sdO, 2 subdwarfs, 1 composite object and 1 emission line star. When we compare this sample with the sources in our UV-excess sample the distribution of spectral types is similar. The fraction of white dwarfs and O/B stars is very different for both surveys, which might be due to the magnitude limit 13$<$$B$$<$16 of the Sandage survey and the criteria used for the classification.\\ \item A second sample of 46 UV-bright sources from the Sandage Two-color Survey obtained by L\'epine et al. (2011) contains 29 DA white dwarfs (63$\%$), 5 DB white dwarfs (11$\%$), 3 DC white dwarfs, 1 DZ white dwarf, 1 DA+dM, 1 sdB, 2 sdO and 4 F-type stars. Here the F-type stars are at Galactic latitudes larger than $|b|>$5. When we compare this sample with the sources in the UV-excess sample the distribution of spectral types and their variety is very similar, e.g. fraction of white dwarfs. The number of H$\alpha$ emission line objects in the Sandage survey is very different from our UV-excess sample.\\ \item The Kitt Peak-Downes (KPD) survey (Downes et al., 1986) sample of 158 UV-excess objects at Galactic latitude $|b|$$<$12\degr, brighter than $B$$<$15.3 and $(U-B)$$<$--0.5 contains 21 DA white dwarfs, 13 white dwarfs of other types, 20 sdO, 40 sdB, 5 Planetary Nebulae, 41 Be stars, 9 Cataclysmic Variables and 9 other peculiar sources. Remarkable is the small fraction of white dwarfs (only $22\%$) and the high number of Be stars and Planetary Nebulae in the KPD survey compared to the \UVEX survey. This might be partly due to the Galactic latitude difference of the two surveys and the fact that the detection of Planetary Nebulae is strongly affected by interstellar obscuration (Fig.\,6 of Miszalski et al., 2008, Fig.\,7 of Parker et al., 2006 and Moe and De Marco, 2006) at Galactic latitudes smaller than $|b|$$<$5. Normally narrow-band and red/IR surveys would be required to select new Planetary Nebulae. The low number of main-sequence sources in the KPD survey can be explained by the demand $(U-B)$$<$--0.5 and their classification of all blue continuum spectra with strong Balmer lines as sdB candidates. This also directly explains why the number of sdO and sdB stars in KPD is reversed compared to \UVEX. Despite the different magnitude depths and colour cuts the fraction of e.g. QSOs, Cataclysmic Variables and DC white dwarfs is the same for both surveys. The distribution of different spectral types over Galactic latitude and Galactic longitude varies strongly as can be seen in Fig.\ \ref{fig:LatitudeVSLongitude}. When we compare only the KPD sources at Galactic latitude smaller than $|b|<$5 and do not take the H$\alpha$ emitters into account, the result is 16 DA, 2 DB, 1 DC, 13 sdB and 4 sdO stars. This result is similar to our classified UV-excess spectra.\\ \end{itemize} | 12 | 6 | 1206.7012 |
1206 | 1206.5199_arXiv.txt | {We present the gamma-ray, X-ray, optical and radio data for GRB100814A. At the end of the slow decline phase of the X-ray and optical afterglow, a sudden and prominent rebrightening in the optical band occurs followed by a fast decay in both bands. This optical rebrightening is accompanied by possible chromatic variations. We discuss possible interpretations, such as double component scenarios and internal dissipation mechanism, with their virtues and drawbacks. We also compare GRB100814A with other {\it Swift} bursts that show optical rebrightenings with similar properties.} \FullConference{Gamma-Ray Bursts 2012 Conference -GRB2012,\\ May 07-11, 2012\\ Munich, Germany} \begin{document} | GRB10814A is an event with differing X-ray and optical behaviour, in particular displaying a late optical rebrightening which may be chromatic. This event joins a sample of bursts showing the same intriguing properties. We have tested three possibilities to explain this behaviour: presence of two jets, combination of reverse and forward shock, and an internal emission mechanism. We find that the first two scenarios require conditions which are unlikely or can account only for a few of the observed features, and the third scenario may give an explanation only by not making clear predictions. | 12 | 6 | 1206.5199 |
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1206 | 1206.0543_arXiv.txt | {\lsi\ is a member of the select group of gamma-ray binaries: galactic binary systems that contain a massive star and a compact object, show a changing milliarcsecond morphology and a similar broad spectral energy distribution (SED) that peaks at MeV-TeV energies and is modulated by the orbital motion. The nature of the compact object is unclear in \lsi , \ls\ and \hess , whereas \psr\ harbours a 47.74\,ms radio pulsar.} {A scenario in which a young pulsar wind interacts with the stellar wind has been proposed to explain the very high energy (VHE, $E$\,$>$\,100\,GeV) gamma-ray emission detected from \lsi , although no pulses have been reported from this system at any wavelength. We aim to find evidence of the pulsar nature of the compact object.} {We performed phased array observations with the Giant Metrewave Radio Telescope (GMRT) at 1280\,MHz centred at phase 0.54. Simultaneous data from the multi-bit phased array (PA) back-end with a sampling time of $t_\mathrm{samp}$\,$=$\,$128\,\mathrm{\mu s}$ and from the polarimeter (PMT) back-end with $t_\mathrm{samp}$\,$=$\,$256\,\mathrm{\mu s}$ where taken.} {No pulses have been found in the data set, with a minimum detectable mean flux density of $\sim$\,0.38\,mJy at 8-$\sigma$ level for the pulsed emission from a putative pulsar with period $P$\,$>$2\,ms and duty cycle $D$\,$=$10\% in the direction of \lsi .} {The detection of posible radio pulsations will require deep and sensitive observations at frequencies $\sim$\,0.5--5\,GHz and orbital phases 0.6$-$0.7. However, it may be unfeasible to detect pulses if the putative pulsar is not beamed at the Earth or if there is a strong absorption within the binary system.} | \label{sec:intro} Located at a distance of $2.0\pm0.2\,\rm{kpc}$ \citep{frail91}, \lsi\ contains a rapidly rotating B0\,Ve star with a stable equatorial shell, and a compact object of unknown nature with a mass between 1 and 5\,$\mathrm{M}_{\sun}$, orbiting it every 26.5\,days \citep{hutchings81,casares05a}. Optical and IR orbital modulation have been found \citep{mendelson94,paredes94}. Using radial velocities data, \citet{casares05a} found that the orbit is eccentric ($e$\,$\simeq$\,0.72) and periastron takes place at phase $0.23\pm0.02$, taking $T_0$\,$=$\,$\rm{JD}\,2,443,366.775$. More recently \citet{aragona09} stated that the eccentricity is slightly lower ($e$\,$\simeq$\,0.54) and the phase of the periastron passage is 0.275. The compact object would be a neutron star for inclinations $25\degr$\,$\la$\,$i$\,$\la$\,$60\degr$ and a black hole if $i$\,$\la$\,$25\degr$ \citep{casares05a}. However, there are discrepancies between the proposed orbital solutions by \citet{casares05a} and \citet{aragona09}, and the inclination angle is poorly constrained by radial velocities of the companion star alone. Quasi-periodic radio outbursts monitored during 23 years have provided an accurate orbital period value of $26.4960\pm0.0028\,\rm{d}$ \citep{gregory02}. The maximum of the radio outbursts varies between phase 0.45 and 0.95. Orbital X-ray periodicity has also been found \citep{paredes97,harrison00,torres10}. Similar results have been obtained at higher energies with \textit{INTEGRAL} data \citep{hermsen06,zhang10}. \lsi\ is also spatially coincident with a high energy (HE, $E$\,$>$\,100\,MeV) gamma-ray source detected by EGRET \citep{kniffen97}. The \textit{Fermi Space Telescope} Large Area Telescope (\textit{Fermi}/LAT) reported the first detection of the orbital modulation at HE with a period of $26.6\pm0.5$\,days roughly anti-correlated with the X-ray emission \citep{abdo09a}. At very high energy (VHE, $E$\,$>$\,100\,GeV) gamma rays, \lsi\ was detected by the MAGIC Cherenkov telescope \citep{albert06} and confirmed by the VERITAS stereoscopic array \citep{acciari09b}. Further observations by the MAGIC collaboration led to the discovery of the orbital variability of TeV emission with a period of $26.8\pm0.2$\,\rm{days} \citep{albert09}. Simultaneous observations of MAGIC and the \textit{XMM-Newton} and \textit{Swift} X-ray satellites revealed a correlation between the X-rays and VHE bands \citep{anderhub09}. \citet{massi04} reported the discovery of an extended jet-like and precessing radio emitting structure at angular extensions of 10--50 milliarcseconds. Owing to the presence of apparently relativistic radio emitting jets, \lsi\ was proposed to be a microquasar. However, VLBA images obtained during a full orbital cycle show a rotating elongated morphology \citep{dhawan06}, which may be consistent with a model based on the interaction between the relativistic wind of a young non-accreting pulsar and the wind of the donor star \citep{maraschi81,dubus06}. On 2008 September 10, the \textit{Swift} Burst Alert Telescope (BAT) detected a hard X-ray burst in the direction of \lsi\ which, assuming the association, would be the signature of a magnetar-like activity that has been proposed to be linked to the presence of a young highly magnetized pulsar in the binary system \citep{dubus08}. The \textit{Fermi}/LAT spectrum is compatible with a power law and an exponential cutoff at $\sim$\,6\,GeV, suggesting that there are two separate spectral components for the HE and VHE emission. Moreover, the spectral similarity with gamma-ray pulsars leads to the consideration of a magnetospheric origin for this HE component despite the fact that no pulses have been detected and although the orbital modulation would be unexpected in this scenario \citep{abdo09a, petri11}. This orbital variability could be understood in the framework of inverse Compton scattering of photon fields in a striped pulsar wind model, which predicts pulsed and variable HE emission \citep{petri11}. In any case, from an observational point of view it is not clear yet if \lsi\ contains an accreting black hole, an accreting neutron star or a non-accreting neutron star. In addition to \lsi , three binary systems that contain a massive star and a compact object and display extended radio emission have been clearly detected at VHE: \psr , \ls\ and \hess . These four systems have a similar spectral energy distribution (SED; \citealt{dubus06, hinton09}), peaking at MeV-GeV energies, and thus are considered gamma-ray binaries. Whereas the nature of the compact object in \lsi , \ls\ and \hess\ remains unknown, \psr /\object{LS~2883} is the only gamma-ray binary with a confirmed pulsar \citep{johnston92}. This binary system contains an O9.5\,Ve \citep{negueruela11} donor and a 47.74\,ms radio pulsar orbiting it every 3.4\,years in a very eccentric orbit with $e$\,$=$\,$0.87$ \citep{johnston92,johnston94}. The radio light curve of the unpulsed emission of \psr\ is well explained by the adiabatic expansion of a synchrotron bubble, which is similar to the behaviour found in \lsi . Pulsed radio emission is detected with a spectral index of about $-0.6$, although the radio pulses vanish for some weeks during the periastron passage \citep{connors02,johnston05} probably due to free-free absorption by the stellar wind and interaction with the Oe-star disk. Note that in the case of \lsi\ the apastron separation is smaller than the periastron separation for \psr , therefore detecting pulses from \lsi\ is unlikely even if the pulsar exists and is beamed at the Earth. The system was detected by HESS in the 2004 and 2007 periastron passages \citep{aharonian05,aharonian09}. Extended and variable radio emission at milliarcsecond scales, outside the binary system, was observed after the 2007 periastron passage of \psr\ \citep{moldon11b}. This is the first observational evidence that a non-accreting pulsar orbiting a massive star can produce a milliarcsecond radio structure similar to those of \lsi , \ls\ and \hess . As in the case of \lsi , the nature of the compact object in \ls\ is unknown: it is either a black hole or a neutron star with a mass between 1.5 and 10\,$\mathrm{M}_{\sun}$ in a slightly eccentric ($e$\,$=$\,0.35) 3.9-day orbit around an O6.5\,V((f)) star located at 2.5\,kpc \citep{casares05b,aragona09}. Similarly to \lsi\ and \psr , \ls\ also shows X-rays, HE and VHE emission modulated by the orbit (\citealt{takahashi09,abdo09b,aharonian06}; respectively). The discovery of a bipolar extended milliarcsecond radio emission morphology with VLBA observations prompted a microquasar interpretation \citep{paredes00}, but additional VLBA observations revealed a changing behaviour of the morphology that cannot easily be explained by a microquasar scenario \citep{ribo08}. \citet{rea11b} performed recently a deep search for pulsations from \ls\ with \textit{Chandra}, finding no periodic signals in a frequency range of 0.005--175\,Hz. \hess\ has recently joined the short list of gamma-ray binaries. \citet{moldon11a} reported the slightly extended and variable radio milliarcsecond structure of \hess\ at 1.6\,GHz with the e-EVN. It is positionally coincident with a B0pe star and the X-ray source XMMU~J063259.3+054801 \citep{hinton09}, which is also variable \citep{acciari09a}. \citet{bongiorno11} have recently confirmed its binary nature by the discovery of a periodicity of 320\,$\pm$\,5\,days in a 0.3--10\,keV light curve obtained with \textit{Swift}. This faint and point-like VHE source was first detected in the HESS Galactic Plane Survey \citep{aharonian07}. \hess\ is clearly variable also at VHE: 2006-2007 and 2008-2009 observations by VERITAS imposed flux upper limits well below the values detected by the HESS Galactic Plane Survey \citep{acciari09a}, and afterwards higher TeV gamma-ray emission was detected by VERITAS \citep{ong11} and MAGIC \citep{mariotti11}. No pulses have been detected at any wavelength from \hess\ \citep{rea11a}. Overall, there are many similarities between \psr\ and the other three known gamma-ray binaries (\lsi , \ls\ and \hess ), but no pulsations have been detected up to now in the latter. In this paper we report our search for radio pulsations from \lsi . In Sect.~\ref{sec:previous} we present the situation of the searches for pulsations from \lsi\ before our work. In Sect.~\ref{sec:absorption} we describe how we selected the orbital phase of the observation and the frequency in order to minimise the effect of absorption and pulse broadening in the observations described in Sect.~\ref{sec:observations}. We describe the data analysis in Sect.~\ref{sec:analysis}. The results are presented in Sect.~\ref{sec:results}, where we describe the implications for future observations. In Sect.~\ref{sec:discussion} we discuss these results and present our conclusions. | \label{sec:discussion} Detection of pulses is the only definitive way to confirm the pulsar nature of the compact object in \lsi . We have performed phased array observations with the GMRT at 1280\,MHz in the \lsi\ direction centred at phase 0.54, from 21:31:45 2009 July 2 UT to 03:43:06 2009 July 3 UT. No pulses have been found in the data set, with a minimum detectable flux density of around 0.38\,mJy for periods greater than 2\,ms and a duty cycle of 10\%. If \lsi\ had the same duty cycle as \psr\ (63\%), the upper limit would be 1.3\,mJy (see Fig.~\ref{fig:ul}), which is below the $\sim$\,5\,mJy mean flux density of this pulsar at this frequency \citep{johnston05}. These results give a guide for future observations planning, which would be best at $\sim$\,0.5--5\,GHz because the upper limit at 1.28\,GHz implies deep upper limits at high frequencies in case of a steep spectrum, as shown in Fig.~4. Following the discussion from Sect.~\ref{sec:absorption}, the orbital phases at which the flux is higher are $\sim$\,0.6$-$0.7, and at these phases the higher flux densities are received at $\sim$\,0.5--2\,GHz assuming a spectral index $\alpha$\,$=$\,$-1.0$. In addition to this, below 0.5\,GHz the pulse broadening due to the ISM scattering is very important for fast pulsars (see Appendix~\ref{app} for more details). Therefore, long ($\sim$\,ten hours) observations at $\sim$\,0.5--5\,GHz with Effelsberg, EVLA or more likely GBT directly pointing to \lsi\ around phases 0.6$-$0.7 are recommended to significantly improve these upper limits or have a good chance to detect the pulses. \citet{mcswain11} performed a pulsar search in \lsi\ using GBT at a similar frequency range (bands C, S and X) and a different orbital phase (0.9), placing deep upper limits (4.1--14.5\,$\mu$Jy). In this argumentation we are not considering the possibility of the pulsar beam not pointing close to our line of sight. This issue is less problematic at X-rays, because the beam is usually wider than the radio one. However, no X-ray pulsations have been detected so far from \lsi . A deep search performed with \textit{Chandra} observations provided $3\sigma$ upper limits on the pulsed fraction of the total X-ray emission ranging between 7 and 15\% in the 0.3$-$10 keV energy band \citep{rea10}. It should be mentioned that the TeV binary system hosting \psr\ does not show X-ray pulsations during its periastron passage \citep{chernyakova06}, with an upper limit of 15\% on the pulsed fraction \citep{rea10}, yet it is known to be a pulsar beamed at Earth. An alternative way to experimentally prove the presence of a pulsar in \lsi\ is finding HE pulsations. As mentioned in Sect.~\ref{sec:intro}, the HE emission detected by \textit{Fermi}/LAT could be explained by inverse Compton scattering of photon fields in a striped pulsar wind model, which predicts pulsed and variable HE emission \citep{petri11}. A better orbital solution would improve the opacity estimation and the knowledge of the acceleration suffered by the putative pulsar and therefore the chances of resolving the pulses in observations along a large part of the orbit. In any case, because of the opacity of the binary system and the possibility of the putative pulsar not beaming towards Earth, it might be impossible to detect pulses even with infinite sensitivity. | 12 | 6 | 1206.0543 |
1206 | 1206.0296_arXiv.txt | Atmospheric mass-loss from Hot-jupiters can be large due to the close proximity of these planets to their host star and the strong radiation the planetary atmosphere receives. On Earth, a major contribution to the acceleration of atmospheric ions comes from the vertical separation of ions and electrons, and the generation of the ambipolar electric field. This process, known as the "polar wind", is responsible for the transport of ionospheric constituents to the Earth's magnetosphere, where they are well observed. The polar wind can also be enhanced by a relatively small fraction of super-thermal electrons (photoelectrons) generated by photoionization. We formulate a simplified calculation of the effect of the ambipolar electric field and the photoelectrons on the ion scale-height in a generalized manner. We find that the ion scale-height can be increased by a factor of 2-15 due to the polar wind effects. We also estimate a lower limit of an order of magnitude increase of the ion density and the atmospheric mass-loss rate when polar wind effects are included. | \label{sec:Intro} In the past two decades and in particular, following the {\it Kepler} mission, hundreds of exoplanets have been detected \citep[e.g.,][]{exoplanet95,exoplanets03}. Many of these planets are gas giants observed at an extremely close orbit of less than 0.1~AU from their host star (an orbital period of less than 10 days), and are classified under the term "Hot-Jupiters" (HJ). The unexpected close-in orbit of HJ has stimulated many science investigations regarding their formation, evolution, and tidal interaction \citep[e.g.,][and references therein]{Papaloizou07}, their magnetic interaction with the host star \citep[e.g.,][and references therein]{Cohen10}, and the structure and dynamics of their atmospheres \citep[e.g.,][and references therein]{Showman08}. In such a close orbit (especially if the star and the planet are tidally-locked), HJ are expected to receive extremely large amounts of stellar X-ray and EUV radiation \citep{Penz08,Cecchi-Pestellini09}. It has been argued that this high EUV radiation can lead to a strong photo-evaporation of the planetary atmosphere and high mass loss rates \citep{Lammer03,Baraffe04, Baraffe06}, leading to a less massive planets. However, this could not be supported by the observed mass distribution \citep{Hubbard07}. Observations of $Ly\alpha$ emission from the HD209458 system have suggested that the planet occupies an inflated hydrogen corona with outflow velocities of $50-100\;km\;s^{-1}$ and a mass-loss rate of about $10^{10}\;g\;s^{-1}$. However, there is a debate on whether the observations are effected by the host star or whether the observed features are of planetary origin \citep{Vidal-Madjar03,Ben-Jaffel07,Vidal-Madjar08}. A more recent observation of the system, as well as of the HD189733 system reveled a smaller mass-loss rate of about $10^{8}\;g\;s^{-1}$ \citep{LecavelierDesEtangs10,Linsky10}. On the theoretical side, several models for atmospheric escape from HJ have been developed in recent years. A detailed models for the chemistry, photoionization, and aeronomy of HJ were developed by \cite{Yelle04} and by \cite{Garcia-Munoz07}. \cite{Tian05} and \cite{Murray-clay09} performed hydrodynamic calculations of thermally driven atmospheric escape, and \cite{Stone09}, \cite{Trammell11}, and \cite{Adams11} included the planetary magnetic field geometry, which confines the escaping gas to regions of open field lines. The models above predict mass-loss rates not higher than $10^{10}\;g\;s^{-1}$. Some of the models also included the incoming stellar wind and found that the planetary outflow is ought to be suppressed by the wind. Non of the models predicts a sufficiently high mass-loss rate so that the planet can be evaporated in a relatively short time-scale. In the Earth's upper atmosphere (as well as in other planets), there is a well observed physical process which plays an important role in the acceleration of ions. The polar wind \citep{BanksHolzer68} is the outflow of planetary ions along open field lines. The main driver for this process is the ambibolar electric field, which is proportional to the electron pressure gradient. Since electrons are more mobile than ions, a charge separation is created along the magnetic field direction, leading to an electric potential that acts on the ions to retain charge neutrality. The end result is an acceleration of the ions by this electric field so that the ions are dragged by the electrons. The electric field applies a force propotional to the negative gradient of the electron pressure. Using some simplifications, the resulting force is approximately equivalent to half of the gravitational force on the major ion species and directed oppositely. Since $O^{+}$ is the major ion species in Earth's upper ionosphere, the result is a supersonic flow of H$^{+}$ and an increase in the O$^{+}$ scale height. In addition, photoelectrons, which are highly energized electrons due to photoionization (the tail of the distribution function), can significantly increase the electron temperature, leading to an enhancement of the ions acceleration \citep{Lemaire72}. In the Earth's upper atmosphere, the velocity of $O^+$ is lower than the escape velocity. Nevertheless, $O^+$ is observed to serve as a significant plasma source in the magnetosphere \citep{Lennartsson81}. \cite{Tam95} and \cite{Tam98} have demonstrated by numerical simulation that photoelectrons indeed, can accelerate $O^+$ and $H^+$, while they obtained an unrealistic electron temperature of $40,000\;K$. An additional simulation by \cite{khazanov97} resulted in a more realistic electron temperature of $16,000\;K$. Recent numerical simulations by \cite{Glocer12} also included the effects of photoelectrons to look at the global outflow solution and compared with in-situ observations. Their simulations showed that the polar wind mechanism is responsible to the transport of ionospheric $H^+$ and $O^+$, and that only a small fraction of photoelectrons can significantly contribute to the ion acceleration. In this letter, we investigate how the ambipolar electric field and the fraction of photoelectrons reduce the gravitational potential, and therefore, increase the ion scale-height and the ion density at the top of the atmosphere of HJ. We also calculate how the mass-loss rate for $H^+$ is modified by this effect. Due to the high EUV radiation, the fraction of photoelectron in the atmospheres of HJ is expected to be higher than in the Earth's case, leading to a much greater increase of electron temperature. In Section~\ref{sec:AEF}, we calculate the change of the effective gravity and the ion scale-height due to the ambipolar electric field and photoelectrons. We present and discuss the results in Section~\ref{sec:Results}, and draw our conclusion in Section~\ref{sec:Conclusions}. | \label{sec:Conclusions} In this paper, we perform a simplified calculation of the effect of the ambipolar electric field and atmospheric photoelectrons on the planetary ion scale-height. We show that this effect can reduce the effective gravity and therefore, enhance the ion acceleration in the region of the planetary atmosphere, where magnetic field lines are open. We find that a small fraction of photoelectrons (less than 1\% of the total electrons) can increase the ion scale-height by a factor of 2-15. We calculate the hydrostatic density profiles using the modified scale-heights and find that the planetary mass-loss rate should increase by an order of magnitude at a minimum, even neglecting any increase in the ion velocity due to this the process. Since the ion acceleration should be enhanced by the process, we expect the increase in mass-loss rate to be even greater. A more comprehensive calculation, however, requires a more detailed modeling effort. | 12 | 6 | 1206.0296 |
1206 | 1206.5016_arXiv.txt | {In this first paper of a series, we report the creation of large and well-defined database that combines extensive new measurements and a literature search of 3876 supernovae (SNe) and their 3679 host galaxies located in the sky area covered by the Sloan Digital Sky Survey (SDSS) Data Release 8 (DR8).} {This database should be much larger than previous ones, and should contain a homogenous set of global parameters of SN hosts, including morphological classifications and measures of nuclear activity.} {The measurements of apparent magnitudes, diameters $(D_{25})$, axial ratios $(b/a)$, and position angles (PA) of SN host galaxies were made using the images extracted from the SDSS $g$-band. For each host galaxy, we analyzed RGB images of the SDSS to accurately measure the position of its nucleus to provide the SDSS name. With these images, we also provide the host galaxy's morphological type, and note if it has a bar, a disturbed disk, and whether it is part of an interacting or merging system. In addition, the SDSS nuclear spectra were analyzed to diagnose the central power source of the galaxies. Special attention was paid to collect accurate data on the spectroscopic classes, coordinates, offsets of SNe, and heliocentric redshifts of the host galaxies.} {Identification of the host galaxy sample is 91\% complete (with 3536 SNe in 3340 hosts), of which the SDSS names of $\sim$1100 anonymous hosts are listed for the first time. The morphological classification is available for 2104 host galaxies, including 73 (56) hosts in interacting (merging) systems. The \emph{total sample} of host galaxies collects heliocentric redshifts for 3317 ($\sim$90\%) galaxies. The $g$-band magnitudes, $D_{25}$, $b/a$, and PA are available for 2030 hosts of the morphologically classified sample of galaxies. Nuclear activity measures are provided for 1189 host galaxies. We analyze and discuss many selection effects and biases that can significantly affect any future analysis of our sample.} {The creation of this large database will help to better understand how the different types of SNe are correlated with the properties of the nuclei and global physical parameters of the host galaxies, and minimize possible selection effects and errors that often arise when data are selected from different sources and catalogs.} | A crucial aspect of many recent studies of extragalactic supernovae (SNe) is to establish the links between the nature of SN progenitors and stellar populations of their host galaxies. The most direct method for realizing this is through their identification on pre-SN images. However, the number of such SNe is small and is limited to the nearby core-collapse (CC) events \citep[e.g.,][]{2009ARA&A..47...63S}. This limitation and small-number statistics are the main reasons to investigate the properties of SN progenitors through indirect methods. In this context, the properties of SN host galaxies, such as the morphology, color, nuclear activity, star formation rate, metallicity, stellar population, age etc. provide strong clues to the understanding of the progenitors \citep[e.g.,][]{1995A&A...297...49P,1997Ap.....40..296K,2002MNRAS.331L..25B, 2003A&A...406..259P,2005A&A...433..807M,2008ApJ...673..999P, 2009A&A...503..137B,2010ApJ...721..777A,2010ApJ...724..502H,2011arXiv1110.1377K}. In addition, valuable information of the nature of progenitors can be obtained through the study of the spatial distribution of SNe \citep[e.g.,][]{2008MNRAS.388L..74F,2009A&A...508.1259H,2010MNRAS.405.2529W} and environments \citep[e.g.,][]{2008MNRAS.390.1527A,2011ApJ...731L...4M,2011A&A...530A..95L}. Over the past decade, many studies have investigated the nature of SN progenitors in the nearby Universe via local and global properties of their host galaxies. For example, \citet{2008ApJ...673..999P} investigated how the different types of SNe are correlated with the metallicity of their host galaxy. They showed strong evidence that SNe~Ibc\footnote{{\tiny We use SNe~Ibc to mainly denote the Ib, Ic and mixed Ib/c SNe types whose specific Ib or Ic classification is uncertain.}} occur in higher metallicity hosts than SNe~II, while there is no such effect for SNe~Ia relative to SNe~II. \citet{2008MNRAS.388L..74F} studied the radial distribution of SNe~Ia in morphologically selected early-type host galaxies from the SDSS, and found that there is no statistically significant difference between the radial distribution of SNe~Ia and the light profile of their early-type host galaxies. \citet{2009A&A...508.1259H} compared the radial distribution of CC SNe within their spiral hosts with the distributions of stars and ionized gas in spiral disks. They concluded that the normalized radial distribution of all CC SNe is consistent with an exponential law, the scale length of the distribution of SNe~II appears to be significantly larger than that of the stellar disks of their host galaxies, while SNe~Ibc have a significantly smaller scale length than SNe~II \citep[see also][]{2009MNRAS.399..559A}. The scale length of the radial distribution of CC~SNe shows no significant correlation with the host galaxy morphological type, or the presence of bar structure. Several authors have studied the radial distributions of SNe of different types in large numbers of galaxies \citep[e.g.,][]{1975PASJ...27..411I,1992A&A...264..428B,1997AJ....113..197V, 2008MNRAS.388L..74F,2009A&A...508.1259H}, but none of these studies attempted to categorize the hosts according to their activity level. However, other authors have shown that the SNe distributions in galaxies with various activity levels might be different \citep[e.g.,][]{1990A&A...239...63P,2005AJ....129.1369P, 2008Ap.....51...69H,2010MNRAS.405.2529W,2012A&A...540L...5H}. For example, \citet{2010MNRAS.405.2529W} directly measured number and surface density distributions of SNe~II in their hosts, and indicated that SNe~II detected in star-forming galaxies follow an exponential law, in contrast, the distribution of SNe~II detected in Active Galactic Nuclei (AGN) hosts significantly deviates from an exponential law. \citet{2005AJ....129.1369P} studying a sample of CC SNe in galaxies hosting AGN found that SNe~Ibc in active/star-forming galaxies are more centrally concentrated than are the SNe~II, but given the small sample, this difference was not statistically significant. The results of \citeauthor{2005AJ....129.1369P} were confirmed with larger samples of CC~SNe by \citet{2008Ap.....51...69H}. The locations of SN explosions in multiple galaxy systems have also been studied. In interacting galaxies, CC~SNe are not preferentially located towards the companion galaxy \citep[e.g.,][]{2001MNRAS.328.1181N}. Similarly, the azimuthal distributions inside the host members of galaxy groups are consistent with being isotropic \citep{2001MNRAS.328.1181N}. \citet{2010ApJ...724..502H} found that SNe~Ia are more likely to occur in isolated galaxies without close neighbors. However, many similar studies \citep[e.g.,][]{2009MNRAS.399..559A,2009A&A...508.1259H,2010MNRAS.405.2529W} presented above are suffered from poor statistics, as well as strong biases on the SNe and their host galaxies sample. Often, they were random selections of nearby SNe and their hosts from the Asiago Supernova Catalogue\footnote{\texttt{{\tiny http://web.oapd.inaf.it/supern/cat/}}} \citep[ASC,][]{1999A&AS..139..531B} or the Sternberg Astronomical Institute (SAI) Supernova Catalogue\footnote{\texttt{{\tiny http://www.sai.msu.su/sn/sncat/}}} \citep[SSC,][]{2004AstL...30..729T} or the official list of all the discovered SNe on the Central Bureau for Astronomical Telegrams (CBAT) website\footnote{\texttt{{\tiny http://www.cbat.eps.harvard.edu/lists/Supernovae.html}}}. Most recently, \citet{2012A&A...538A.120L} published a unified SN catalog for 5526 extragalactic SN that were discovered up to 2010 December~31. The unified catalog mostly combines ASC, SSC, and data from CBAT in a consistent way, and adopts all of the inhomogeneous data on SNe and their host galaxies from the original sources. For the galaxy data, these catalogs made large use of the Third Reference Catalogue of Bright Galaxies (RC3) by \citet{1991rc3..book.....D} and the HyperLeda\footnote{\texttt{{\tiny http://leda.univ-lyon1.fr/}}} database \citep{2003A&A...412...45P} as well as the NASA/IPAC Extragalactic Database\footnote{\texttt{{\tiny http://ned.ipac.caltech.edu/}}} (NED). Hence, the data are given with various degrees of accuracy depending on the accuracy of the original catalog. Many selection effects and errors that often arise when data is selected from different sources and catalogs can significantly bias results and lead to wrong conclusions. Quantitative studies of SN progenitors therefore require a large and well-defined sample of SNe and their host galaxies, and our goal is to provide the database for such a sample. The Eighth Data Release\footnote{\texttt{{\tiny http://www.sdss3.org/}}} (DR8) of the SDSS \citep{2011ApJS..193...29A} covering over 14000 square degrees with high quality imaging and spectroscopy makes it finally possible to construct better samples for studies of the properties of the host galaxies and environments of SNe. This large amount of the SDSS data enables statistically meaningful studies that are only little affected by selection effects. In this paper, we report the creation of a large database of several thousand SNe that exploded in galaxies identified in the SDSS DR8. We provide identifications of SN host galaxies, their accurate coordinates, heliocentric redshifts, morphological types, and activity classes, as well as apparent magnitudes, diameters, axial ratios, and position angles. However, our goal is not just to increase the size of the sample in comparison with previous studies, but also to carry out morphological classification, as well as individual measurements of the global parameters of the host galaxies in a homogenous way. In addition, we summarize the overall statistical properties of the sample, analyzing and discussing residual selection effects and biases that can still affect subsequent studies and results. An additional motivation for this study is the comparison between the ASC, SSC, and CBAT databases, to reveal possible inconsistencies in the listed SNe types and offsets. Our database includes corrected data with their uncertainties on the SNe. This is the first paper of a series and is organized as follows: Sect.~\ref{data} introduces the data and describes in detail the reduction techniques. In Sect.~\ref{Resdiscus}, we give the results and discuss all the statistical properties and selection effects of the sample. A summary and perspectives for future uses are finally addressed in Sect.~\ref{sumandper}. Throughout this paper, we adopt a cosmological model with $\Omega_{\rm m}=0.27$, $\Omega_{\rm \Lambda}=0.73$, and the Hubble constant of $H_0=73 \,\rm km \,s^{-1} \,Mpc^{-1}$ \citep{2007ApJS..170..377S}, consistent with direct determination based on Cepheid variables and SNe~Ia by \citet{2009ApJ...699..539R}. Future papers of this series will use this database to determine how the different types of SN progenitors are correlated with the global parameters (morphology, size, luminosities etc.) of the host galaxies as well as on their nuclear properties (activity class, metallicity, stellar population etc.). | \label{Resdiscus} In this section, we compare some of our measured parameters with previously available measurements, provide distributions and statistics on several of them, study their evolution with distance, and finally discuss various selection effects. This analysis is important to understand whether our sample of SN and their host galaxies from the SDSS DR8 is representative of the general population. Our comparisons with measurements from HyperLeda and SDSS are done by performing robust linear regressions, with iterative rejection of $3\,\sigma$ outliers, and we also measure a robust estimate of the dispersion, using the median absolute deviation (MAD) of the residuals, converting to $\sigma_{\rm MAD} = 1.483\,\rm MAD$, where the numerical factor is the one appropriate for Gaussian distributions of residuals. \subsection{Comparison of host galaxy morphological classifications} It is well known that the central parts of the images of many very bright galaxies and high surface brightness objects may be over-exposed, which affects their morphological classification. Meanwhile, galaxies that are faint (low surface brightness) can also be misclassified, because of the lack of precise morphological details in the image. Using a homogenous sample of 604 SNe, \cite{2002PASP..114..820V,2003PASP..115.1280V,2005PASP..117..773V} classified the SN host galaxies from the Lick Observatory Supernova Search (LOSS) in the David Dunlap Observatory (DDO) morphological type system. They suggested that to understand the dependence of SN type on the host galaxy population, it is more important to obtain accurate morphological classifications than it is to increase the size of the sample. For example, among $\sim$800 morphologically classified hosts of CC~SNe, \cite{2008A&A...488..523H} found 22 cases where the host had been classified as E or S0. Following a detailed morphological analysis, they found that among these 22 early-type objects, 17 are in fact misclassified spiral galaxies, one is a misclassified irregular, and one is a misclassified ring galaxy, leaving only 3 early-type galaxies\footnote{ One of these 3 SNe, SN~2005md, reported by \cite{2005CBET..332....1L} and initially classified by \cite{2005IAUC.8650....2M} as a probable young type~IIb SN, was shown to be in fact a new Galactic cataclysmic variable \citep{2010ATel.2750....1L}.}. In this respect, the host morphology is a crucial parameter in the study of SN progenitors. To present a detailed numerical comparison of our morphological classification of SN host galaxies with those given in the SN catalogs (mainly from RC3), we introduce our $t$-type values that we will use in this study. In comparison to the standard RC3 classification, we have grouped Ellipticals and Lenticulars into broader classes: cE and E galaxies are typed together, as well as S0$^{-}$, S0$^{0}$ and S0$^{+}$ galaxies. Indeed, using the SDSS images, it was not possible to visually distinguish the differences between these subclasses. Table~\ref{ttype} shows the relation between our $t$-types and those of the RC3. In the top panel of Fig.~\ref{morphRC3}, we show the comparison of our morphological classifications for the 1313 hosts from our classified sample that are also present in the HyperLeda database (when both classifications are more accurate than just ``S'': 1767 for our sample among the total of 2104 classified galaxies). The reference system of the morphological classification in HyperLeda is generally the RC3. The comparisons with RC3 types were performed after converting the RC3 numerical ($t$) classifications to our scheme of $t$ versus type (see Table~\ref{ttype}). Point sizes in the figure correspond to the number of hosts in each morphological bin. In the bottom panel of Fig.~\ref{morphRC3}, we present the distributions of the differences between the HyperLeda $t$-types and ours. \begin{table}[t] \begin{center} \caption{Relation between the RC3 morphological types and $t$ values with ours. \label{ttype}} \begin{tabular}{lrclr} \hline \hline \multicolumn{2}{c}{Ours} & & \multicolumn{2}{c}{RC3} \\ \cline{1-2} \cline{4-5} \multicolumn{1}{c}{Type} & \multicolumn{1}{c}{$t$} & & \multicolumn{1}{c}{Type} & \multicolumn{1}{c}{$t$} \\ \hline E & --3 & & cE & --6 \\ E & --3 & & E & --5 \\ E/S0 & --2 & & E$^{+}$ & --4 \\ S0 & --1 & & S0$^{-}$ & --3 \\ S0 & --1 & & S0$^{0}$ & --2 \\ S0 & --1 & & S0$^{+}$ & --1 \\ S0/a & 0 & & S0/a & 0 \\ Sa & 1 & & Sa & 1 \\ Sab & 2 & & Sab & 2 \\ Sb & 3 & & Sb & 3 \\ Sbc & 4 & & Sbc & 4 \\ Sc & 5 & & Sc & 5 \\ Scd & 6 & & Scd & 6 \\ Sd & 7 & & Sd & 7 \\ Sdm & 8 & & Sdm & 8 \\ Sm & 9 & & Sm & 9 \\ Im & 10 & & Im & 10 \\ \hline \end{tabular} \end{center} \end{table} \begin{figure}[h] \begin{center} \includegraphics[width=0.9\hsize]{morph.ps} \end{center} \caption{\emph{Top}: comparison of HyperLeda (reordered) versus our $t$ morphological types for 1313 host galaxies with available classifications. Point sizes are keyed to the number of objects. \emph{Bottom}: the distributions of differences of our and HyperLeda $t$-types according to our classification. The error bars for the mean values in each bin are presented. The \emph{solid lines} in each figure are added to visually better illustrate the deviation in classifications.} \label{morphRC3} \end{figure} \begin{figure*}[ht] \begin{center}$ \begin{array}{rlrl} \includegraphics[width=0.24\hsize]{Fig5.ps} & \includegraphics[width=0.24\hsize]{Fig6.ps} & \includegraphics[width=0.24\hsize]{Fig7.ps} & \includegraphics[width=0.24\hsize]{Fig8.ps} \end{array}$ \end{center} \caption{Examples of the SDSS DR8 and DSS~I images of SN host galaxies representing the cases when $| t_{\rm our}-t_{\rm HyperLeda} | \geq 3$. The PGC object identifier is listed at the top with our (\emph{left}) and the HyperLeda (mainly selected from the RC3, \emph{right}) classifications. In all images, north is up and east to the left.} \label{fmorph} \end{figure*} Inspection of Fig.~\ref{morphRC3} shows a trend for our classifications to be later overall (except E) in comparison with those of HyperLeda. The mean deviation in classifications is $0.65 \pm 0.04$ $t$-types. Meanwhile, the mean absolute deviation is $1.05 \pm 0.03$ $t$-types. Table~\ref{tab:compare} shows the results of a linear regression between the HyperLeda types and ours. The relation between the two measures of $t$ has a slope significantly different than unity, with a best fit value of 0.93, and the residuals from this trend have a dispersion close to 1.5 types. A similar trend was already found by \cite{2010ApJS..186..427N}, who recently released a morphological catalog of 14034 visually classified galaxies $(0.01 < z < 0.1)$. \begin{table}[t] \begin{center} \caption{Comparisons of HyperLeda and SDSS~DR7 measurements with ours for SN host galaxies. \label{tab:compare}} \tabcolsep 4pt \begin{tabular}{llrrr@{\hspace{5.5mm}}} \hline \hline Quantity & Ref. cat. & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} & \multicolumn{1}{c}{Dispersion} \\ \hline $t$ & HyperLeda & $0.93\pm0.01$ & $-0.29\pm0.05$ & 1.48 \\ $\log a$ & HyperLeda & $0.93\pm0.00$ & $0.06\pm0.01$ & 0.06 \\ $\log a$ & SDSS & $0.87\pm0.01$ & $0.25\pm0.01$ & 0.07 \\ $b/a$ & HyperLeda & $0.94\pm0.01$ & $0.04\pm0.01$ & 0.08 \\ $b/a$ & SDSS & $0.97\pm0.01$ & $0.01\pm0.01$ & 0.07 \\ PA & HyperLeda & $1.00\pm0.00$ & $-0.94\pm0.49$ & 5.72 \\ PA & SDSS & $1.00\pm0.00$ & $-0.56\pm0.47$ & 4.57 \\ mag & HyperLeda & $0.98\pm0.00$ & $0.66\pm0.06$ & 0.22 \\ mag & SDSS & $0.87\pm0.00$ & $2.01\pm0.06$ & 0.25 \\ \hline \\ \end{tabular} \parbox{\hsize}{\textbf{Notes.} Columns~3 ($a$) and 4 ($b$) represent the robust linear fits (with iterative rejection of outliers) of $x_{\rm theirs} = a\,x_{\rm ours}+b$. The last column (Dispersion) is computed as 1.483 times the median absolute deviation (which corresponds to $\sigma$ for Gaussian distributions) of the residuals from our best-fit trend. Cases with no measurements (given arbitrary values such as 99) have been discarded. The SDSS magnitude is cModelMag in the $g$-band. The PA dispersion is in degrees.} \end{center} \end{table} \begin{figure}[t] \begin{center} \includegraphics[width=0.9\hsize]{morph_bar.ps} \end{center} \caption{Distributions of morphological types and presence of bars (\emph{shaded region}) in 1767 classified hosts.} \label{morphbar} \end{figure} We found that approximately 10\% (125 hosts) of the 1313 galaxies in common with HyperLeda have $t$-types that are dramatically different $(|t_{\rm our}-t_{\rm HyperLeda}| \geq 3)$. In Fig.~\ref{fmorph}, we present several extreme cases where the difference between morphological type codes is $\geq 3$. Color images are taken from the SDSS DR8 on which our classification was performed, while the grayscale images are from the photographic plates given in the Digitized Sky Survey~I (DSS~I), as the RC3 classification is mostly based on these or similar plates. In many cases, photographic plates suffer from the narrow dynamical range that causes saturation as well as underexposure, and also from their non-linear response functions \citep[e.g.,][]{1995MNRAS.274.1107N}. A detailed comparative study of the SDSS and DSS images of hosts, when $|t_{\rm our}-t_{\rm HyperLeda}| \geq 3$, allows us to emphasize that in nearly all cases, the overexposure as well as low resolution of the photographic plates cause late-type galaxies of high surface brightness to be misclassified as early-type in the RC3. We have found a handful of cases with the opposite trend: E/S0 misclassified as spirals in RC3, mainly due to the heterogeneous nature of morphological data sets in the HyperLeda. \cite{2011A&A...532A..74B} have recently released the EFIGI (Extraction de Formes Id\'{e}alis\'{e}es de Galaxies en Imagerie) catalogue; a multi-wavelength database specifically designed to densely sample all Hubble types. Their imaging data were obtained from the SDSS DR4 for a sample of 4458 PGC galaxies. This catalog includes 453 galaxies from our sample. We found very good agreement between the EFIGI morphological classifications and ours. \begin{figure}[t] \begin{center} \includegraphics[width=0.95\hsize]{type_d.ps} \end{center} \caption{\emph{Top}: distribution of morphological types as a function of distance. The types of SN host galaxies have been grouped into the following broad classes: E and E/S0 (\emph{gray}), S0 and S0/a (\emph{red}), Sa and Sab (\emph{magenta}), Sb and Sbc (\emph{orange}), Sc and Scd (\emph{green}), and hosts Sd to Im (\emph{blue}). \emph{Bottom}: fractional distribution of morphological types (with the same morphological groups) as a function of distance.} \label{dmorph} \end{figure} Fig.~\ref{morphbar} shows the distribution (see Table~\ref{Gmorph}) of SN host galaxies with respect to our $t$-type. This histogram shows that the intermediate types Sb, Sbc, and Sc are the most frequent. Only Sm and Im types, which are intrinsically faint, have fewer than 30 galaxies per type. In the top panel of Fig.~\ref{dmorph} we present the histogram of morphological types as a function of distance. The types of host galaxies have been grouped into the following broad classes: E and E/S0 (gray), S0 and S0/a (red), Sa and Sab (magenta), Sb and Sbc (orange), Sc and Scd (green), and hosts Sd to Im (blue). The bottom panel of Fig.~\ref{dmorph} shows the fractional distribution of types of the same morphological groups as a function of distance. Host galaxies with $D < 200~{\rm Mpc}$ include nearly the whole range in morphological types. At the same time, the late-type hosts are preferentially distributed in the lower distance bins, while early-types are more populated at the higher distances. At distances $D < 200~{\rm Mpc}$, Sd to Im type hosts represent 12\% of the classified galaxies. At large distances, spiral galaxies can be under-represented, because spiral arms are difficult to resolve \citep[see, e.g.,][]{2010ApJS..186..427N,2011A&A...532A..74B}. Hence, the most massive/luminous early-type galaxies prefer the higher distances, while the least massive/luminous late-type galaxies are more populated at the lower distances. This is also the selection effect on the SN type: bright SNe Ia, exploding also in E galaxies, can be found more easily than fainter CC~SNe, exploding only in late-type galaxies \citep[e.g.,][]{2011MNRAS.412.1419L}. Therefore, the sample of CC~SNe hosts is closer on average than the sample of SNe Ia host galaxies. \subsection{Comparison of presence of bars in host galaxies} A proper detection of barred structures of hosts is very important when constraining the nature of the SNe progenitors by comparing their distribution within host galaxies with the distributions of stellar populations and ionized gas in the disks \citep[e.g.,][]{2005AJ....129.1369P,2009A&A...508.1259H}. Below, we have carried out a comparative study to find differences in the detection of barred structures of host galaxies between our classified sample and HyperLeda. Fig.~\ref{morphbar} presents the distribution (see Table~\ref{Gmorph}) of hosts with or without bars as a function of $t$-type. Roughly 29\% of our 1767 classified galaxies have bars. The barred fraction is highest in types Sb, Sbc, and Sc. A detailed comparison with HyperLeda reveals that in 378 galaxies among our 1767 in common (21\%), HyperLeda fails to detect the bar that we visually detect on the SDSS images or conversely detects a bar when we don't. In Fig.~\ref{bar}, we present examples of hosts galaxies with discrepancies in bar detection between HyperLeda and us. Given their superior angular resolution and 3-colour representations, the SDSS images offer a much more reliable source for bar detection than do the plate-based images on which most of the HyperLeda classifications were performed. Inspecting these cases of discrepancies in bar detection, we conclude that HyperLeda fails to show bars in both high central surface brightness early-type galaxies and late-type galaxies with low surface brightness bars. We found that bars tend to be incorrectly detected in HyperLeda galaxies of high inclination ($i>70^\circ$). The remaining cases of detection discrepancies are again due to the heterogeneous nature of the HyperLeda data sets. Note that we may also have missed weak bars because of inclination effects, or that in some cases the SDSS images of hosts may be too shallow to detect bars. For instance, among our S0-Sm galaxies with inclinations $i<70^\circ$, the average bar fraction is $(37\pm1)\%$ whereas for hosts with inclinations $\geq70^\circ$ the average bar fraction is only $(11\pm2)\%$. \subsection{Comparison of isophotal measurements of host galaxies} \begin{figure*} \begin{center}$ \begin{array}{rlrl} \includegraphics[width=0.24\hsize]{Fig1.ps} & \includegraphics[width=0.24\hsize]{Fig2.ps} & \includegraphics[width=0.24\hsize]{Fig3.ps} & \includegraphics[width=0.24\hsize]{Fig4.ps} \end{array}$ \end{center} \caption{Examples of SDSS DR8 and DSS~I images of SN host galaxies representing the cases of discrepancies in bar detection between HyperLeda and us. The PGC object identifier is listed at the top with our (\emph{left}) and the HyperLeda (mainly selected from the RC3, \emph{right}) classifications. In the two galaxies on the \emph{left}, we detect bars, while HyperLeda does not, while in the two galaxies on the \emph{right}, HyperLeda claims that there is a bar while we dot not see any on our higher-resolution 3-band images. In all images, north is up and east to the left.} \label{bar} \end{figure*} We checked the differences of our measurements of $g$-band major axes with the HyperLeda $B$-band diameters $(D_{25})$ as well as with the SDSS $g$-band isophotal major axes (isoA). Table~\ref{tab:compare} shows the results of a linear regression between the HyperLeda, SDSS isophotal measurements and ours. In the top left panels of Figs.~\ref{galourHL} and \ref{galourSDSS} we show the discrepancy of the major diameters of the host galaxies between the samples. Our measured $g$-band diameters are systematically larger than the HyperLeda $D_{25}$ diameters in $B$-band for all the morphological types of hosts. Our diameters are greater than those in the HyperLeda on average by a factor of $1.32\pm0.01$. This level of discrepancy is not unexpected, given that our sizes are measured at the $\mu_g=25\,\rm mag\,arcsec^{-2}$ isophotal level. With the transformation equation, $B=g+0.39\,(g-r)+0.21$ (\citealp{2005AJ....130..873J}, for all stars with $R-I<1.15)$, and given our mean $g-r\simeq 0.7$ color, our $g$-band measurements are performed at the equivalent of the $\left\langle\mu_B\right\rangle\simeq 25.48$ isophote, hence our greater host galaxy sizes. For the data comparisons with SDSS, we used DR7 instead of DR8, because DR7 includes isophotal photometric quantities. Before comparing the samples of classified E-Im galaxies we excluded 96 objects with unreliable SDSS measurements that differ from our measurements by more than a factor of 2. We considered the data for these galaxies as incorrectly measured in SDSS. Our diameters are in good agreement with isoA, and greater only on average by a factor of $1.01\pm0.01$. In general, the SDSS measurements are unreliable for objects larger than $\sim$100~arcsec \citep[e.g.,][]{2011A&A...532A..74B}. Still, for objects smaller than $100''$, we find scatters corresponding to factors of 33\% (with HyperLeda) and 16\% (with SDSS~DR7). In the top right panels of Figs.~\ref{galourHL} and \ref{galourSDSS} we present the differences of our measured axial ratios and that of the HyperLeda as well as isoB/isoA of the SDSS DR7. The majority ($\sim$94\%) of E-Im galaxies, for which the axial ratios are available in our and other measurements, these axial ratios are consistent within 0.2. The mean deviation of our measurements from that of the HyperLeda is $0.005\pm0.003$, and from isoB/isoA is $0.015\pm0.003$. The MADs are 0.055 with HyperLeda and 0.052 with SDSS~DR7. In fact, after correction for trends, the residuals show a robust dispersion of $\sim$0.1. There is no dependence of residuals on the morphological types of host galaxies. The position angles (PAs) of the major axes were determined at the same (${\mu_g = \rm 25\,\ mag\,\ arcsec^{-2}}$) isophotal level as the measurements of angular diameters. Comparisons of our PA measurements with the HyperLeda and SDSS DR7 (isoPhi) determinations are shown in the bottom left panels of Figs.~\ref{galourHL} and~\ref{galourSDSS}. There are non-negligible fractions of cases where $\rm PA_{\rm our} + PA_{\rm theirs} = 180^\circ$, especially for $\rm PA_{\rm our} \simeq 0^\circ$ or $180^\circ$, for which small errors can (for the correct sign) flip the PA to $180^\circ$ minus its true value. To avoid these unfair extreme outliers in our comparisons, we redistributed the values of differences of PAs such that we considered the $\Delta$PA$-$180$^\circ$ or $\Delta$PA$+$180$^\circ$ when the differences were $> 90^\circ$ or $< -90^\circ$, respectively. \begin{figure*}[t] \begin{center} \includegraphics[width=0.85\hsize]{our_hyp.ps} \end{center} \caption{\emph{Top left}: comparison between our measurements of major axes of E-Im galaxies and those of the HyperLeda in the $B$-band. \emph{Top right}: comparison between measured axial ratios and those of the HyperLeda. \emph{Bottom left}: comparison between our measurements of position angles and those of the HyperLeda. \emph{Bottom right}: comparison between our measurements of apparent $g$-band magnitudes and that in the $B$-band of the HyperLeda. The color coding corresponds to Fig.~\ref{dmorph}: E-E/S0 (\emph{gray filled circles}), S0-S0/a (\emph{red open circles}), Sa-Sab (\emph{magenta triangles}), Sb-Sbc (\emph{orange crosses}), Sc-Scd (\emph{green filled squares}), and Sd-Im (\emph{blue open squares}). The \emph{solid lines} in each figure are added to visually better illustrate the deviations. The \emph{dashed lines} are best fit linear trends from Table~\ref{tab:compare}. \label{galourHL}} \end{figure*} \begin{figure*}[t] \begin{center} \includegraphics[width=0.85\hsize]{our_sdss.ps} \end{center} \caption{\emph{Top left}: comparison between our measurements of major axes of E-Im galaxies and isophotal major axes in the $g$-band of the SDSS DR7. \emph{Top right}: comparison between measured axial ratios and those of the SDSS DR7. \emph{Bottom left}: comparison between our measurements of position angles and those of the SDSS DR7. \emph{Bottom right}: comparison between our measurements of apparent $g$-band magnitudes and Composite Model Magnitudes of the same band in the SDSS DR7. The color and symbol coding corresponds to Fig.~\ref{galourHL}. The \emph{solid lines} in each figure are added to visually better illustrate the deviations. The \emph{dashed lines} are best fit linear trends from Table~\ref{tab:compare}. \label{galourSDSS}} \end{figure*} With these corrections, the mean difference between HyperLeda's PAs and ours is $0\fdg9\pm0\fdg5$, while the MAD is $4\fdg9$. The comparison with the PAs from SDSS DR7 (isoPhi) yields a mean deviation of $0\fdg5\pm0\fdg6$ and a MAD of $4\fdg3$. In both cases, 85\% of the host galaxies have PAs consistent within $20^\circ$ with those of HyperLeda or SDSS. The scatter in the bottom left panel of Fig.~\ref{galourHL} may be due to the fact that the HyperLeda values correspond to measurements made at $\mu_{\rm B}={\rm 25\,\ mag\,\ arcsec^{-2}}$, whereas ours are made at typically lower surface brightness thresholds: $\langle\mu_{\rm B}\rangle=$$\,\ \sim$${\rm 25.48\,\ mag\,\ arcsec^{-2}}$. Inspections of the remaining cases with large discrepancies show that they are mostly contributed by peculiar and low surface brightness galaxies or objects in interacting/merging systems. In addition, the PA is hard to determine when the galaxy is face-on, because for any given elliptical aperture it is ill-defined; 97\% of galaxies with $b/a\leq0.5$ are consistent within $20^\circ$ of PA, whereas only 80\% of hosts with $b/a>0.5$ have the same consistency. Again, there is no dependence of discrepancies on the morphological types of galaxies. For more detailed inspection and explanation of this effect see \cite{2007ApJS..170...33P}. We also compared our $g$-band magnitude measurements to the HyperLeda $B$-band and to the SDSS DR7 Composite Model Magnitude (cModelMag) determinations in the $g$-band, which are measured from the linear combination of the exponential and de~Vaucouleurs profiles that fit best the $g$-band SDSS images. Also, there is excellent agreement between cModelMag and \cite{1976ApJ...209L...1P} magnitudes of galaxies. Although, the cModelMag and Petrosian magnitudes are not identical, there is an offset of 0.05-0.1 mag but this is within errors of our elliptical aperture measurements. The results of magnitude comparisons are presented in the bottom right panels of Figs.~\ref{galourHL} and \ref{galourSDSS}. The relation between the HyperLeda magnitudes and ours is quite linear with a slope of 0.98 (our magnitudes are slightly brighter relative to theirs at the bright end). The residuals from our robust linear fit between the two magnitudes is 0.22 mag. The mean difference between our magnitudes and those of HyperLeda is $-0.42\pm0.01$, while the mean absolute difference is $0.47\pm0.01$. It is clear that most of the measurements agree well with each other once the magnitudes are brought into the same system. For $\sim$65\% of galaxies our and HyperLeda magnitude differences are less than 0.5~mag, which corresponds to 0.02~mag after converting the $g$-band into $B$-band. Only for $\sim$7\% of the hosts galaxies the magnitude difference is equal or larger than 1~mag, hence 0.52~mag after the conversion. These galaxies are mostly of types Sc to Im. The mean difference between our magnitudes and the SDSS cModelMag photometry is $-0.21\pm0.01$. The MAD is 0.19 mag. There is a trend where we measure brighter magnitudes than SDSS for the brightest galaxies. Despite some curvature in the relation between the two magnitude estimates, we fit a lines, and find $m_{\rm SDSS} = 0.87\,m_{\rm ours}$, with a dispersion of 0.25 mag about this relation. For $\sim$81\% of galaxies our and the SDSS magnitude differences are less than 0.5~mag. For $\sim$5\% of the objects the magnitude differences are equal to or larger than 1~mag. For faint galaxies, the SDSS measurements algorithm overestimates the cModelMag fluxes, while for bright galaxies it underestimates the fluxes. This trend is stronger for late-type galaxies. Again, we explain the presence of the large scatters, especially for the bright galaxies, by the unreliable SDSS photometric measurements for objects larger than $\sim$100~arcsec \citep[e.g.,][]{2011A&A...532A..74B}. We also performed the same analysis using the cModelMag of SDSS DR8 instead of DR7, and found the same behavior for the photometric bias. \begin{table*}[t] \begin{center} \caption{Distribution of SN types according to the morphological classification of the host galaxies.} \label{bigtable} \begin{tabular}{lrrrrrrrrrrrrrrrrr} \hline \hline &\multicolumn{1}{c}{E}&\multicolumn{1}{c}{E/S0}&\multicolumn{1}{c}{S0}&\multicolumn{1}{c}{S0/a}&\multicolumn{1}{c}{Sa} &\multicolumn{1}{c}{Sab}&\multicolumn{1}{c}{Sb}&\multicolumn{1}{c}{Sbc}&\multicolumn{1}{c}{Sc}&\multicolumn{1}{c}{Scd} &\multicolumn{1}{c}{Sd}&\multicolumn{1}{c}{Sdm}&\multicolumn{1}{c}{Sm}&\multicolumn{1}{c}{Im}&\multicolumn{1}{c}{S} &\multicolumn{1}{c}{Unclassified}&\multicolumn{1}{r}{All}\\ \hline I&4&2&3&5&4&4&8&6&8&1&1&1&3&0&4&18&72\\ Ia&61&35&51&73&31&50&122&125&114&28&22&9&2&7&199&1061&1990\\ Ib&1&0&1&0&1&2&6&9&14&5&4&3&2&1&5&9&63\\ Ib/c&0&0&0&1&3&2&5&12&10&1&2&3&0&0&4&8&51\\ Ic&0&0&0&0&1&4&15&30&18&9&10&1&1&0&8&23&120\\ II&0&0&2&4&9&9&98&121&182&39&39&15&9&12&55&116&710\\ IIb&0&0&0&0&2&0&4&9&10&5&3&5&2&1&6&3&50\\ IIn&0&0&0&1&2&0&11&12&28&9&3&3&3&2&9&27&110\\ Unclassified&9&9&18&16&17&21&68&69&65&14&17&10&6&4&55&312&710\\ \hline All&75&46&75&100&70&92&337&393&449&111&101&50&28&27&345&1577&3876\\ \hline \\ \end{tabular} \parbox{\hsize}{\textbf{Notes.} All SNe types include uncertain (``:'' or ``?'') and peculiar (``pec'') classifications. Type~II SNe include subtypes II~P and II~L. Types~I, Ia, and II include also few SNe classified from the light curve only, these SNe are labeled by ``*'' symbols in the \emph{total sample}.} \end{center} \end{table*} \begin{figure}[t] \begin{center} \includegraphics[width=\hsize]{plot.ps} \end{center} \caption{\emph{Top left}: distribution of the inclination angles for disk (S0-Sm) galaxies. The \emph{dashed} curve represents the expected random distribution. \emph{Top right}: distribution of corrected $g$-band absolute magnitudes for the classified (E-Im) galaxies. The \emph{dashed} histogram shows the distribution of $g$-band absolute magnitudes of the SDSS Main Galaxy sample (the values are divided by 500 for the sake of clarity). \emph{Bottom}: the $g$-band luminosity of the same galaxies as a function of their distance. The color and symbol coding corresponds to Fig.~\ref{galourHL}. The average luminosity in different distance bins is overplotted as big open circles with error bars of the mean values. The dashed line represents the selection limit of the SDSS Main Galaxy spectroscopic sample ($r\leq 17.77$) for the extinction corrected Petrosian magnitude, assuming $g-r=0.64$.} \label{galaxyfdata} \end{figure} In addition, we checked the influence of nuclear activity on the discrepancies of the photometric measurements. Indeed, since the SDSS total magnitudes for extended objects (ModelMag and cModelMag) are based on single-component fits, we could expect that such fitting will perform poorly for AGNs with relatively bright nuclei. The result is negative: there is no dependence of the discrepancies on the nuclear activity of the host galaxies. The top left panel of Fig.~\ref{galaxyfdata} shows the distribution of inclination angles for morphologically classified disk (S0-Sm) galaxies. There is a clear deficit of SN host galaxies having small and large inclinations. We thus share the view with \cite{2011MNRAS.412.1419L}, who found a deficit of LOSS galaxies with small inclinations and explained this deficit by limits of the precision on the major and minor axes. Indeed, it is very difficult to measure inclinations smaller than $20^\circ$ from elliptical aperture measurements applied to nearly face-on galaxies. The lack of galaxies with large inclinations can be explained by a bias in the discovery of SNe (see also the middle panel of Fig.~\ref{dSNtype}) in highly inclined spirals \citep[e.g.,][]{1988A&A...190...10C,1991ARA&A..29..363V,1997A&A...322..431C}. \begin{figure}[t] \begin{center} \includegraphics[width=\hsize]{sntype_d.ps} \end{center} \caption{\emph{Left}: distribution of different types of SNe as a function of the host galaxy $t$-type. \emph{Middle}: distribution of SNe as a function of the inclination of the spiral hosts. \emph{Right}: distribution of different types of SNe as a function of the distance of their host galaxy.} \label{dSNtype} \end{figure} The top right panel of Fig.~\ref{galaxyfdata} presents the distribution of corrected $g$-band absolute magnitudes for the classified galaxies. In comparison with the distribution of SDSS galaxies (dashed histogram), the SN host galaxies are more luminous. A Kolmogorov-Smirnov (KS) test indicates that the more luminous distribution of SN host galaxy magnitudes, relative to the SDSS galaxies in general, cannot be obtained by chance with more than 0.1\% probability. This distribution again suffers from a selection effect on SN productivity, since the rate of SNe depends on the luminosity or stellar content of the host galaxies \citep[e.g.,][]{1991ARA&A..29..363V,1997A&A...322..431C, 2011MNRAS.412.1473L,2011Ap.....54..301H}. Therefore, our sample of classified host galaxies is biased toward bright galaxies. For comparison, the distribution of $g$-band absolute magnitudes of the SDSS Main Galaxy sample is also shown. The bottom panel of Fig.~\ref{galaxyfdata} shows the $g$-band luminosity of the same galaxies as a function of their distance. Galaxy luminosities were derived from absolute magnitudes, assuming that $g$-band absolute magnitude of the Sun is 5.45 \citep{2003ApJ...592..819B}. The luminosities of late-type (Sd-Im) hosts are on average 5 times lower than those of early-type (E-E/S0) galaxies. This trend is clearly seen in the bottom panel of Fig.~\ref{galaxyfdata}. The average $g$-band luminosities in different distance bins are also plotted. At greater distances, the low-luminosity host galaxies are lost due to flux limitations. Thus, Malmquist bias causes the average $g$-band luminosity to increase with increasing distance. The databases become progressively incomplete for low-luminosity galaxies at greater distances. This was already mentioned by \cite{2011MNRAS.412.1419L} for the LOSS galaxy sample. \begin{figure}[t] \begin{center}$ \begin{array}{cc} \includegraphics[width=0.48\hsize]{Fig11.ps} & \includegraphics[width=0.48\hsize]{Fig12.ps} \end{array}$ \end{center} \caption{SDSS images of early-type (E or S0) host galaxies with CC~SNe. The objects identifiers are listed at the top with our morphological classification. The SN names, types, and positions (marked by cross sign) are also shown. In all images, north is up and east to the left.} \label{CCinES0} \end{figure} \subsection{Distribution of SN types} In Table~\ref{bigtable}, we present the distribution of SN types according to the morphology of their host galaxies in an analogous fashion as in Table~5 of \cite{1999A&AS..139..531B}. It is clear that most SNe are found in spiral hosts. The percentage of unclassified SNe is $\sim$18\% of the \emph{total sample}. Approximately half of the unclassified SNe were discovered before 2000. The left panel of Fig.~\ref{dSNtype} shows the number distribution of SN types as a function of host morphology. There is a significant difference between the distribution of the SN~Ia hosts and that of CC SN hosts, while types~Ibc and II SN hosts have similar distributions. These trends were previously reported by \cite{2002PASP..114..820V,2003PASP..115.1280V, 2005PASP..117..773V} and \cite{2011MNRAS.412.1419L}. \begin{table}[t] \begin{center} \caption{Distribution of SN types according to the level of nuclear activity of the host galaxy. \label{actSNe}} \begin{tabular}{llrrrrr} \hline \hline Diagram & Activity & Ia & Ibc & II & Unclassified & All \\ \hline & Sy & 20 & 3 & 14 & 4 & 41 \\ & LINER & 77 & 8 & 41 & 40 & 166 \\ BPT& C & 78 & 12 & 64 & 27 & 181 \\ & SF & 110 & 49 & 157 & 96 & 412 \\ \cline{2-7} & All & 285 & 72 & 276 & 167 & 800 \\ \hline & Sy & 83 & 7 & 43 & 35 & 168 \\ & LINER & 38 & 7 & 24 & 19 & 88 \\ WHAN& SF & 199 & 66 & 218 & 121 & 604 \\ & RP & 235 & 13 & 72 & 86 & 406 \\ \cline{2-7} & All & 555 & 93 & 357 & 261 & 1266 \\ \hline \\ \end{tabular} \parbox{\hsize}{\textbf{Notes.} Each of the Seyfert (Sy) rows additionally include 6 BL~AGN hosts. The RP row includes 77 P galaxies. The column of unclassified SNe includes also type I SNe.} \end{center} \end{table} It is generally believed that hosts of CC~SNe are objects with young stellar populations (generally spiral or irregular galaxies), while the old stellar population of early-type galaxies can produce only SNe Ia \citep[e.g.,][]{1991ARA&A..29..363V,1999A&A...351..459C}. Nevertheless, among the morphologically classified host galaxies of CC~SNe in our sample, we have found 4 cases (2000ds [Ib] in PGC~025915, 2006ee [II] in PGC~007536, 2007ke [Ib] in PGC~010959, and 2009fe [II] in SBS~1646+499) in which the host has been classified as E or S0, in apparent contradiction to this conventional view. Fig.~\ref{CCinES0} presents the cases of CC~SNe in early-type hosts. \cite{2008A&A...488..523H} already reported and investigated in detail two cases of such CC~SNe in early-type galaxies (2000ds and 2006ee). The host galaxy of SN~2000ds \citep{2000IAUC.7511....2F} has been confirmed to be an S0, with a central region showing dust and a disky central gas distribution \citep[e.g.,][]{2008A&A...488..523H}. According to the outer isophotal structure and radial surface brightness profile of the host of SN~2006ee \citep{2006IAUC.8741....1P}, this must be an S0 galaxy. It has been shown that the surface brightness distribution has some small degree of asymmetry in the region to the south-southwest of the nucleus \citep{2008A&A...488..523H}. Here, we suspected the presence of an embryonic spiral arm. We classified the host galaxy of SN~2007ke \citep{2007CBET.1101....1F} as type E, this classification is also given in both NED and HyperLeda. It is in interacting system and is a member of the cluster of galaxies. The host galaxy of SN~2009fe \citep{2009CBET.1819....1K} is classified as an uncertain type S0. The same morphological type is reported in \cite{2011Ap.....54...15G}. In NED, it is classified as a blazar (Seyfert~1). This object also shows 1.4 GHz radio continuum emission \citep{1998AJ....115.1693C}. However, more detailed inspection on high resolution images is still required. In principle, the galaxy could have some diffuse spiral arms and be classified as S0/a but due to insufficient resolution of the SDSS image at the distance of this object, it has been classified as an uncertain type S0. The presence of CC~SNe in early-type galaxies can be interpreted as an additional indication that residual star formation episodes also take place in E or S0 galaxies, due to merging/accretion or interaction with close neighbors. Recently, \cite{2009MNRAS.394.1713K} have found that the recent star formation is likely to be driven by minor mergers, which seems to fit with our interpretation as well. Meanwhile, using rest-frame UV photometry of early-type galaxies in the nearby Universe, \cite{2007ApJS..173..619K} suggested that low-level recent star formation is widespread in nearby early-type galaxies. The situation is also very similar at intermediate redshifts \citep{2008MNRAS.388...67K}. Hence, the detection of SNe II in early-type galaxies is expected, but at lower frequency than type~Ia. The middle panel of Fig.~\ref{dSNtype} shows the number distribution of different types of SNe as a function of the inclination of S0-Sm hosts (see also the top left panel of Fig.~\ref{galaxyfdata} and its explanation). The KS test suggests that there is no significant difference between the distributions of types~Ia and CC~SNe. The same behavior occurs when comparing type~Ia SNe separately with types~Ib and II SNe. A similar trend was mentioned by \cite{1997A&A...322..431C}. The right panel of Fig.~\ref{dSNtype} shows the number distribution of SN types as a function of distance. All the major types of SNe are peaked at 50-100 Mpc. It is clear that the sample of SNe is largely incomplete beyond $\sim$100 Mpc. The distributions of types~Ibc and II~SNe are similar and display a sharp decline beyond 100 Mpc. Type Ia SNe, because of their comparatively high luminosity and the presence of dedicated surveys (e.g., SDSS SN Survey, ESSENCE etc.), are discovered at much greater distances than CC~SNe. A similar behavior was also found by \cite{2011MNRAS.412.1419L}. Table~\ref{actSNe} displays the numbers of different types of SNe in hosts with different levels of nuclear activity. It is important to note that nuclear activity is affected by aperture bias \citep[e.g.,][]{2005PASP..117..227K}. The SDSS spectra are taken with a fixed fiber size $(3'')$. For a nearby galaxy, the SDSS fiber covers the central nuclear region or its part, while for more distant case it covers a larger fraction of the galaxy (e.g., 120 pc at $z=0.004$ but 2.7 kpc at $z=0.1$). Hence, the activity data can be contaminated by the emission of circumnuclear regions or from the whole galaxy. The effect also depends on galaxy size, as for dwarf galaxies the fiber will cover a larger fraction of the total emission. | 12 | 6 | 1206.5016 |
1206 | 1206.6709_arXiv.txt | {We present a new observational campaign, DWARF, aimed at detection of circumbinary extrasolar planets using the timing of the minima of low-mass eclipsing binaries. The observations will be performed within an extensive network of relatively small to medium-size telescopes with apertures of $\sim$~20~--~200~cm. The starting sample of the objects to be monitored contains (i) low-mass eclipsing binaries with M and K components, (ii) short-period binaries with sdB or sdO component, and (iii) post-common-envelope systems containing a WD, which enable to determine minima with high precision. Since the amplitude of the timing signal increases with the orbital period of an invisible third component, the timescale of project is long, at least 5-10 years. The paper gives simple formulas to estimate suitability of individual eclipsing binaries for the circumbinary planet detection. Intrinsic variability of the binaries (photospheric spots, flares, pulsation etc.) limiting the accuracy of the minima timing is also discussed. The manuscript also describes the best observing strategy and methods to detect cyclic timing variability in the minima times indicating presence of circumbinary planets. First test observation of the selected targets are presented. } | With the continuing discovery of extrasolar planets and an expectation that the majority of solar-type stars reside in binary or multiple systems (Duquennoy \& Mayor 1991), planetary formation in binary systems has become an important issue (Lee et al. 2009). In general, we can consider planetary companions to binary stars in two types of hierarchical planet-binary configurations: first "S-type" planets which orbit just one of the stars, with the binary period being much longer than that of the planet; second, "P-type" or circumbinary planets, where the planet simultaneously orbits both stars, and the planetary orbital period is much longer than that of the binary (Muterspaugh et al. 2007). Simulations show either of above possibilities has a large range of stable configurations (see e.g., Broucke 2001; Pilat-Lohinger \& Dvorak 2002; Pilat-Lohinger et al. 2003; Benest 2003). Recent theoretical studies (e.g., Moriwaki \& Nakagawa 2004; Quintana \& Lissauer 2006; Pierens \& Nelson 2008) have predicted that P-type planets can form and survive over long timescales. Characterization of such planets is potentially of great interest because they can lead to a better understanding of the formation and evolution of planetary systems around close binary stars, which can be rather different from the case of single stars. (Lee et al. 2009). Hereafter we will consider the "P-type" or circumbinary planets only. The detection of circumbinary planets is far from being easy. Three principal techniques are: (i) precise radial-velocity (hereafter RV) measurements to detect the wobble of the binary mass center (Konacki et al. 2009), (ii) photometric detection of transits of the planet(s) across the disks of the components of the inner binary (Doyle et al. 2011), and (iii) timing of the inner binary eclipses (Lee et al. 2009). The classical RV technique is complicated by the fact that the large RV changes of the binary components mask the small wobble resulting from the circumbinary planet(s). In the case of close binaries, the situation is exacerbated by the tidal spin-up of the components, where the projected rotational velocity ($v \sin i$) often exceeds 100 km~s$^{-1}$ (see the DDO close-binary project, Pribulla et al., 2009 and references therein). This makes the RV precision insufficient to detect any systemic velocity changes. In fact, there are hardly any binaries where the systemic-velocity changes revealed a third component (unseen in spectra). The second technique - to detect circumbinary planets searching for transits of a substellar companion across a close binary - requires very long photometric runs with excellent accuracy. Assuming that the orbital planes of the underlying binary and the outer orbit of the substellar body are close to being coplanar, the method should be advantageous for edge-on eclipsing binaries (EBs). Three such systems were found in the Kepler satellite photometry: Kepler-16b (Doyle et al., 2011), Kepler-34b and Kepler-35b (Welsh et al. 2012). The latter two systems were actually identified by eclipse timing. Even if the substellar component is not transiting the inner binary, it causes timing variations of eclipses of the binary system due to the finite velocity of light (light-time effect, hereafter LITE). The eclipses act as an accurate clock for detecting other objects revolving around the inner binary and to determine their orbital parameters from the $observed-calculated$ times of minima (O-C) in way similar to the solution of RV curves. The timing technique proved to be the most fruitful in detecting circumbinary planets (see Section~\ref{timing_searches}). The principal goal of the project DWARF is to detect circumbinary planets and/or substellar companions through the timing analysis of selected close eclipsing binaries. In this paper, we first summarize previous and/or ongoing searches for timing variability (Section~\ref{timing_searches}) and then describe the target selection criteria for the DWARF project (Section~\ref{sample}). In Section~\ref{reduction} we present the CCD data reduction technique and determination of minima times. The following section describes the modeling of the (O-C) residuals for the observed targets. The telescope network put together for continuous photometric monitoring and follow-up observations is described in Section~\ref{network}. | The presented project is aimed at the detection of circumbinary extrasolar planets and brown dwarfs using minima timing variability of carefully selected EBs. Unlike more widespread techniques (RV or transit searches) to detect extrasolar planets, the minima timing does not require high-end and costly astronomical instrumentation. Precise photometric observations of the brightest targets of our sample can be performed by well-equipped amateur astronomers. The chances to detect circumbinary bodies does not depend only on the precision of the individual timings but also on the number of participating institutions and devoted amateurs and number of targets monitored. The theoretical estimates show that the timing technique enables to detect circumbinary planets down to Jupiter mass orbiting on a few-year orbits. The merit of an EB strongly depends on its brightness, depth, and width of the minima (see Eqn.~\ref{precision_min}), less on the mass of the underlying EB (Eqn.~\ref{lite_amp}). The observations within the project promise additional useful science such as: (i) the study of spot cycles in the RS~CVn-like late-type binaries, detection of flares (see Pribulla et al., 2001), (ii) a more accurate characterization of recently-discovered detached eclipsing binaries, (iii) detection of new low-mass EBs which is crucial to better define the empirical lower main sequence, (iv) determination of absolute parameters of the components (in the case that spectroscopic orbits are available), (v) detection of EBs with pulsating component(s), (vi) detection and characterization of multiple systems with two systems of eclipses, (vii) detection of new variable stars in the CCD fields covered, (viii) photometric detection of transits of substellar components across the disks of the components of the eclipsing pair (see Doyle et al., 2011), (ix) detection of invisible massive components causing precession of the EB orbit and changes of the minima depth (see Mayer et al., 2004). The LITE can always be regarded {\it only} as very good indication of a substellar body in the system. In nearby systems with a sufficiently close visual companion (e.g., CU~Cnc, GK~Boo) the LITE on a long-period orbit could be possibly checked by the differential astrometry of the visual pair. | 12 | 6 | 1206.6709 |
1206 | 1206.1635_arXiv.txt | We present the discovery of KELT-1b, the first transiting low-mass companion from the wide-field Kilodegree Extremely Little Telescope-North (KELT-North) transit survey, which surveys $\sim 40\%$ of the northern sky to search for transiting planets around bright stars. The initial transit signal was robustly identified in the KELT-North survey data, and the low-mass nature of the occultor was confirmed via a combination of follow-up photometry, high-resolution spectroscopy, and radial velocity measurements. False positives are disfavored by the achromaticity of the primary transits in several bands, a lack of evidence for a secondary eclipse, and insignificant bisector variations. A joint analysis of the spectroscopic, radial velocity, and photometric data indicates that the $V=10.7$ primary is a mildly evolved mid-F star with $\teff=6518\pm 50~{\rm K}$, $\log{g_*}=4.229_{-0.019}^{+0.012}$ and $\feh=0.008\pm0.073$, with an inferred mass $M_*=1.324\pm0.026~M_\odot$ and radius $R_*=1.462_{-0.024}^{+0.037}~R_\odot$. The companion is a low-mass brown dwarf or a super-massive planet with mass $M_P=27.23_{-0.48}^{+0.50}~\mjup$, radius $R_P=1.110_{-0.022}^{+0.032}~\rjup$, surface gravity $\log{g_{P}}=4.738_{-0.023}^{+0.017}$, and a density $\rho_P=24.7_{-1.9}^{+1.4}~{\rm g~cm^{-3}}$. The companion is on a very short ($\sim 29$ hour) period circular orbit, with an ephemeris $T_c({\bjdtdb})=2455909.292797\pm 0.00024$ and $P=1.2175007\pm 0.000018~{\rm d}$, and a semimajor axis of $a=0.02466\pm 0.00016$AU. KELT-1b receives a large amount of stellar insolation, with $\langle F \rangle = 7.81_{-0.33}^{+0.42} \times 10^9~{\rm erg~s^{-1}~cm^{-2}}$, implying an equilibrium temperature assuming zero albedo and perfect redistribution of $T_{eq}=2422_{-26}^{+32}~{\rm K}$. Upper limits on the secondary eclipse depth in $i$ and $z$ bands indicate that either the companion must have a non-zero albedo, or it must experience some energy redistribution. Comparison with standard evolutionary models for brown dwarfs suggests that the radius of KELT-1b is likely to be significantly inflated. Adaptive optics imaging reveals a candidate stellar companion to KELT-1 with a separation of $588\pm 1$mas, which is consistent with an M dwarf if it is at the same distance as the primary. Rossiter-McLaughlin measurements during transit imply a projected spin-orbit alignment angle $\lambda = 2\pm16$ degrees, consistent with the orbit pole of KELT-1b being aligned with the spin axis of the primary. Finally, the $\vsini=55.4\pm 2.0~{\rm km~s^{-1}}$ of the primary is consistent at $\sim 2~\sigma$ with tidal synchronization. Given the extreme parameters of the KELT-1 system, we expect it to provide an important testbed for theories of the emplacement and evolution of short-period companions, as well as theories of tidal dissipation and irradiated brown dwarf atmospheres. | The most information-rich exoplanetary systems are those in which the companion happens to transit in front of its parent star. Transiting systems are enormously useful for enabling detailed measurements of a seemingly endless array of physical properties of extrasolar planets and their host stars (see reviews by \citealt{winn2009,winn2010}). The most basic properties that can be measured using transiting planets are the planet mass and radius, and so average density. These parameters alone allow for interesting constraints on the internal composition and structure of planets \citep{guillot05,fortney2007,rogers2010,miller2011}. In addition to these basic parameters, transiting planets enable the study of their atmospheres \citep{seager2000,charbonneau02,vidal03,seagerd2010} and thermal emission \citep{deming05,charbonneau2005,knutson2008}. They also allow measurement of planetary and stellar oblateness, rotation rate, and spin-orbit alignment \citep{seager02,spiegel2007,carter2010,rossiter1924,mclaughlin1924,winn2005,gaudi06}. Transiting planets may also be searched for associated rings and moons \citep{brown01,barnes04,tusnski2011}. Further, variations in transit timing may indicate the presence of other bodies in the system \citep{holman05,agol05,steffen05,ford06,ford07,kipping2009}. With sufficiently precise observations, one may constrain the presence of planets with masses smaller than that of the Earth \citep{agol06,carter2010}. The high scientific value of transiting planet systems motivated the first dedicated wide-field transit surveys, which by now have identified over 100 transiting systems (TrES, \citealt{alonso2004}; XO, \citealt{mccullough2006}; HATNet, \citealt{bakos2007}; SuperWASP, \citealt{cameron2007a}, QES, \citealt{alsubai2011}). Although there is substantial diversity in their design, strategy, and sensitivity, these surveys can be grossly characterized as having relatively small cameras with apertures of order $10$~cm, and detectors with relatively wide fields-of-view of tens of square degrees. These surveys are primarily sensitive to giant, close-in planets with radii $R_P\ga 0.5\rjup$ and periods of $P\la 10~{\rm d}$, orbiting relatively bright FGK stars with $V\sim 10-12$. The space-based missions CoRoT \citep{baglin2003} and Kepler \citep{borucki2010} have dramatically expanded the parameter space of transit surveys, enabling the detection of planets with sizes down to that of the Earth and below, planets with periods of several years, and planets orbiting a much broader range of host stars. Furthermore, their large target samples have allowed the detection of rare and therefore interesting planetary systems. These missions have already announced over 50 confirmed planets, and the Kepler mission has announced an additional $\sim 2300$ candidates \citep{batalha2012}, most of which are smaller than Neptune. Notable individual discoveries include the first detection of a transiting Super-Earth \citep{leger2009}, the detection of a `temperate' gas giant with a relatively long period of $ \sim 100$ days \citep{deeg2010}, the first multi-planet transiting systems \citep{steffen2010,holman2010,latham2011,lissauer2011}, the first circumbinary planets \citep{doyle2011,welsh2012}, and the detection of planets with radius of $\la R_\oplus$ \citep{muirhead2012, fressin2012}. Although Kepler and CoRoT have revolutionized our understanding of the demographics of planets, the opportunities for follow-up of the systems detected by these missions are limited. By design, both missions primarily monitor relatively faint stars with $V\ga 12$. Consequently, many of the follow-up observations discussed above that are generically enabled by transiting systems are not feasible for the systems detected by Kepler and CoRoT. Detailed characterization of the majority of these systems will therefore be difficult or impossible. There is thus an ongoing need to discover transiting planets orbiting the bright stars, as well as to increase the diversity of such systems. All else being equal, the brightest stars hosting transiting planets are the most valuable. Larger photon flux permits more instruments and/or facilities to be employed for follow-up, allows subtler effects to be probed, reduces statistical uncertainties, and generally allows for improved or more extensive calibration procedures that help to control systematic errors. Furthermore, brighter stars are also easier to characterize, and are more likely to have pre-existing information, such as proper motions, parallaxes, metallicities, effective temperatures, angular diameters, and broadband colors. The majority of the brightest ($V\la 8$) FGK dwarfs in the sky have been monitored using precision radial velocity surveys for many years, and as a result most of the giant planets with periods of less than a few years orbiting these stars have already been discovered (e.g., \citealt{wright2012}). A smaller subset of these stars have been monitored over a shorter time baseline with the sensitivity needed to detect Neptune- and SuperEarth-mass planets. Because of the low a priori transit probability for all but short period planets, the transiting systems constitute a very small fraction of this sample. To date, seven planets first discovered via radial velocity have subsequently been discovered to also transit; all of the host stars for these planets are brighter than $V=9$. Although there are projects that aim to increase this sample \citep{kane2009}, the overall yield is expected to be small. Because RV surveys generically require spectroscopic observations that are observationally expensive and must be obtained in series, it is more efficient to discover transiting planets around the much more abundant fainter stars by first searching for the photometric transit signal, and then following these up with targeted RV observations to eliminate false positives and measure the planet mass. However, in order to compensate for the rarity and low duty cycle, many stars must be monitored over a long time baseline. Photometric transit surveys that target brighter stars therefore require larger fields of view. Most of the original transit surveys had fields of view and exposure times that were optimized to detect planets orbiting stars with $V\ga 10$. Indeed, only $\sim 20$ transiting planets orbiting stars with $V\le 10$ are currently known ($\sim 40$ with $V\la 11$). Of those with $V\le 10$, $\sim 40\%$ were originally detected by RV surveys. The Kilodegree Extremely Little Telescope-North (KELT-North) transit survey \citep{KELT_SYNOPTIC} was designed to detect giant, short-period transiting planets orbiting the brightest stars that are not readily accessible to RV surveys. \citet{KELT_THEORY} determined the optimal hardware setup specifically to detect transiting planets orbiting stars with $V\sim 8-10$, and based on the specified design requirements in that paper, the KELT-North survey telescope system was constructed using off-the-shelf, high-end consumer equipment. In fact, as the current detection demonstrates, KELT has exceeded its design goals, and is sensitive to transiting systems in some favorable cases down to $V\sim 12$. In addition to the goal of filling in the magnitude gap between radial velocity and other transit surveys, the KELT-North survey also has the potential to detect fainter systems with $V\ga 10$ that are in the magnitude range of previous surveys, but were missed or overlooked for various reasons. The detection discussed in this paper is an example of this opportunity. Here the fact that the KELT-North survey is only now starting to vet candidates, more than eight years after the first candidates were announced by other transit surveys, can be seen an advantage. In particular, previous surveys have established the existence of massive brown dwarf companions \citep{deleuil2008,irwin2010,bouchy2011a,johnson2011,bouchy2011b}, and have demonstrated the feasibility of detecting low-mass companions to hot, rapidly rotating stars \citep{cameron2010}. Partially in response to these results, the KELT-North survey deliberately broadened our search targets to include hot and/or rapidly-rotating stars, which were previously neglected by many transit surveys. The evolving perception of what kinds of stars constitute viable transit search targets played an interesting role in the discovery of KELT-1b, as discussed in \S\ref{sec:hat}. The KELT-North survey has been collecting data since September 2006, and has acquired a sufficient number of high-quality images to detect transit candidates. We have been systematically identifying and vetting transit candidates since February 2011, and in this paper we report our first confirmed low-mass transiting companion, which we designate KELT-1b. KELT-1b has a mass of $\sim 27~\mjup$, and we will therefore follow convention and refer to it as a `brown dwarf' throughout the majority of this paper. However, as we discuss in \S\ref{sec:bdd}, we are, in fact, agnostic about its true nature and therefore how it should be categorized. The outline of this paper is as follows. In order to introduce the survey and provide the appropriate context for our discovery, in \S\ref{sec:survey} we summarize the properties of the KELT-North survey and our procedure for candidate selection. In \S\ref{sec:observations} we review the observations of KELT-1, starting with the properties of the candidate in the KELT-North data, and then summarize the follow-up photometry, spectroscopy, and high-contrast imaging. \S\ref{sec:char} describes our analysis and characterization of the host star and its substellar companion. In \S\ref{sec:discussion} we provide a speculative discussion of the possible implications of this unique system for theories of the emplacement and tidal evolution of short-period substellar companions, models of the structure and atmosphere of brown dwarfs, and the demographics of substellar companions to stars. We briefly summarize in \S\ref{sec:summary}. | \label{sec:discussion} From our global fit to the light curves and RVs, we find that KELT-1b is a low-mass companion with a measured mass $M_P=27.23_{-0.48}^{+0.50}~\mjup$ and radius $R_P=1.110_{-0.022}^{+0.032}~\rjup$. It is on a circular orbit with a semimajor axis of $a=0.02466\pm 0.00016$AU. The host KELT-1 is a mildly evolved mid-F star with a mass $M_*=1.324\pm0.026~M_\odot$, radius $R_*=1.462_{-0.024}^{+0.037}~R_\odot$, effective temperature $\teff=6518\pm 50~{\rm K}$, and a likely age of $\sim 1.5-2$Gyr. Because of its small semimajor axis and hot host, KELT-1b receives a large stellar insolation flux of $\langle F \rangle = 7.81_{-0.33}^{+0.42} \times 10^9~{\rm erg~s^{-1}~cm^{-2}}$, implying a high equilibrium temperature of $T_{eq}=2422_{-26}^{+32}~{\rm K}$ assuming zero albedo and perfect redistribution. Both the surface gravity and density of KELT-1b are substantially higher than that of its host star, and higher than we would expect for a stellar object. We find that the orbit normal of KELT-1b is well-aligned with the projected rotation axis of its host star, with a projected alignment angle of $\lambda = 2 \pm 16$ degrees. Even among the large and diverse menagerie of known transiting exoplanets and low-mass companions, KELT-1b is unique. First, it is one of only 7 unambiguous objects with mass the range $\sim 13-80~\mjup$ that are known to transit stars. Among these, it has the shortest period, and orbits the brightest host star ($V=10.7$). In addition, there is potentially a stellar M dwarf companion to the primary. For all these reasons, KELT-1b is likely to be a very interesting object for further study, and we expect it will provide a benchmark system to test theories of the emplacement and evolution of short period companions, as well the physics of tidal dissipation and irradiated atmospheres of substellar objects. We will discuss some of these ideas briefly. \subsection{Brown Dwarf or Supermassive Planet? KELT-1b and the Brown Dwarf Desert}\label{sec:bdd} Is KELT-1b a brown dwarf (BD) or a is it a suppermassive planet? By IAU convention, brown dwarfs (BDs) are defined to have masses between the deuterium burning limit of $\sim 13~\mjup$ \citep{spiegel2011} and the hydrogen burning limit of $\sim 80~\mjup$ (e.g., \citealt{burrows1997}). Less massive objects are defined to be planets, whereas more massive objects are stars. By this definition, KELT-1b is a low-mass BD. However, it is interesting to ask whether or not KELT-1b could have plausibly formed in a protoplanetary disk, and therefore might be more appropriate considered a ``supermassive planet'' \citep{schneider2011}. More generally, it is interesting to consider what KELT-1b and systems like it may tell us about the formation mechanisms of close companions with masses in the range of $10-100\mjup$. One of the first results to emerge from precision Doppler searches for exoplanets is the existence of a BD desert, an apparent paucity brown dwarf companions to FGK stars with periods less than a few years, relative to the frequency of stellar companions in the same range of periods \citep{marcy2000}. Subsequent studies uncovered planetary companions to such stars in this range of periods in abundance \citep{cumming08}, indicating that the BD desert is a local nadir in the mass function of companions to FGK stars. The simplest interpretation is that this is the gap between the largest objects that can form in a protoplanetary disk, and the smallest objects that can directly collapse or fragment to form a self-gravitating object in the vicinity of a more massive protostar. Therefore, the location of KELT-1b with respect to the minimum of the brown dwarf mass function might plausibly provide a clue to its origin. \subsubsection{Comparison Sample of Transiting Exoplanets, Brown Dwarfs, and Low-mass Stellar Companions}\label{sec:comp} In order to place the parameters of KELT-1b in context, we construct a sample of transiting exoplanets, BDs, and low-mass stellar companions to main sequence stars. We focus only on transiting objects, which have the advantage that both the mass and radius of the companions are precisely known\footnote{In contrast, for companions detected only via RVs, only the minimum mass is known. Of course, one can make an estimate of the posterior probability distribution of the true mass given a measured minimum mass by adopting a prior for the distribution of inclinations (e.g., \citealt{lee2011}). However, this procedure can be particularly misleading in the case of BDs: if BDs are indeed very rare, then objects with minimum mass in the BD desert are more likely to be stellar companions seen at low inclination. Anecdotally, in those cases where constraints on the inclinations can be made, companions with minimum mass near the middle of the brown dwarf desert often do turn out to be stars (e.g., \citealt{sahlmann2011,fleming2012}).}. We collect the transiting exoplanet systems from the Exoplanet Data Explorer (exoplanets.org, \citealt{wright2011}), discarding systems for which the planet mass is not measured. We supplement this list with known transiting brown dwarfs \citep{deleuil2008,johnson2011,bouchy2011a,bouchy2011b,anderson2011}. We do not include the system discovered by \citep{irwin2010}, because a radius measurement for the brown dwarf was not possible. We also do not include 2M0535$-$05 \citep{stassun2007}, because it is a young, double BD system. We add several transiting low-mass stars near the hydrogen burning limit \citep{pont2005,pont2006b,beatty2007}. We adopt the mass of XO-3b from the discovery paper \citep{johnskrull2008}, which is $M_P=13.1 \pm 0.4~\mjup$, which straddles the deuterium burning limit \citep{spiegel2011}. However, later estimates revised the mass significantly lower to $M_P=11.8\pm 0.6$ \citep{winn2008}. We will therefore categorize XO-3b as an exoplanet. The disadvantage of using samples culled from transit surveys is that the sample size is much smaller, and transit surveys have large and generally unquantified selection biases (e.g., \citealt{gaudi2005,fressin2009}), particularly ground-based transit surveys. We emphasize that such biases are almost certainly present in the sample we construct. We have therefore made no effort to be complete. The comparisons and suggestions we make based on this sample should not be considered definitive, but rather suggestive. Figure \ref{fig:mpvs} places KELT-1b among the demographics of known transiting companions to main sequence stars, focusing on massive exoplanets, BDs, and low-mass stars with short periods of $\la 30$ days. KELT-1b has the tenth shortest period of any transiting exoplanet or BD known. It has the sixth shortest period among giant ($M_P\ga 0.1~\mjup$) planets, with only WASP-19b, WASP-43b, WASP-18b, WASP-12b, OGLE-TR-56b, and HAT-P-23b having shorter periods. KELT-1b is more massive by a factor $\sim 3$ than the most massive of these, WASP-18b \citep{hellier2009}. KELT-1b has a significantly shorter period than any of the previously known transiting brown dwarfs, by a factor of $\ga 3$. KELT-1b therefore appears to be located in a heretofore relatively unpopulated region of the $M_P-P$ parameter space for transiting companions. Although the KELT-1 system is relatively unique, it is worth asking if there are any other known systems that bear some resemblance to it. The $M_P\simeq 18 \mjup$, $P\simeq 1.3$ day RV-discovered companion to the M dwarf HD 41004B \citep{zucker2003} has similar minimum mass and orbit as KELT-1b, however the host star is obviously quite different. Considering the host star properties as well, perhaps the closest analogs are WASP-18b \citep{hellier2009}, WASP-33b \citep{cameron2010}, and KOI-13b \citep{mazeh2012,mislis2012}. All three of these systems consist of relatively massive ($M_p\ga 3~\mjup$) planets in short ($\la 2$ day) orbits around hot $(\teff \ga 6500$~K$)$ stars. The mass of KELT-1b ($\sim 27~\mjup$) is close to the most arid part of the BD desert, estimated to be at a mass of $31_{-18}^{+25}~\mjup$ according to \citet{grether2006}. Thus, under the assumption that the BD desert reflects the difficulty of forming objects with this mass close to the parent star under {\it any} formation scenario, KELT-1b may provide an interesting case to test these various models. For disk scenarios, gravitational instability can likely form such massive objects, but likely only at much larger distances \citep{rafikov2005,dodson2009,kratter2010}. The maximum mass possible from core accretion is poorly understood, but may be as large as $\sim 40~\mjup$ \citep{mordasini2009}. The possibility of significant migration of KELT-1b from its birth location to its present position must also be considered, particularly given the existence of a possible stellar companion to KELT-1 (\S\ref{sec:keckao}). This possibility complicates the interpretation of the formation of KELT-1b significantly. For example, it has been suggested that brown dwarf companions are more common at larger separations \citep{metchev2009}; thus KELT-1b may have formed by collapse or fragmentation at a large separation, and subsequently migrated to its current position via the Kozai-Lidov mechanism \citep{kozai1962,lidov1962}. One clue to the origin of KELT-1b and the BD desert may be found by studying the frequency of close BD companions to stars as a function of the stellar mass or temperature. Figure \ref{fig:mpvs} shows the mass of known transiting short period companions as a function of the effective temperature of the host stars. As pointed out by \citet{bouchy2011b}, companions with $M_P\ga 5~\mjup$ appear to be preferentially found around hot stars with $\teff \ga 6000~{\rm K}$, and KELT-1b follows this trend. Although these hot stars are somewhat more massive, the most dramatic difference between stars hotter and cooler than 6000~K is the depth of their convection zones. This led \citet{bouchy2011b} to suggest that tides may play an important role in shaping the frequency and distribution of massive exoplanet and brown dwarf companions to old stars. Some evidence for this has been reported by \citet{winnrm2010}, who argue that hot ($\teff \ge 6250$K) stars with close companions preferentially have high obliquities, suggesting that if the emplacement mechanisms are similar for all stars, tidal forces must later serve to preferentially bring cool host stars into alignment. Figure \ref{fig:rmdist} shows the distribution of spin-orbit alignments for transiting planets versus the host star effective temperature. KELT-1b falls in the group of hot stars with {\it small} obliquities. Interestingly the other massive $\ga 5~\mjup$ planets are also located in this group. We discuss the possible formation and evolutionary history of KELT-1b, and the likely role of tides in this history, in more detail below. We remain agnostic about the classification of KELT-1b as a brown dwarf or supermassive planet. \begin{figure} \epsscale{1.2} \plotone{rm.eps} \caption{ The projected spin-orbit alignment angle $\lambda$ for transiting planets as measured by the RM effect, versus the effective temperature of the host star, following \citet{winnrm2010}. The grey points show exoplanets with mass $M_P< 5~\mjup$, whereas the black points show those with $M_P> 5 \mjup$. KELT-1b, shown with a star, is the first transiting brown dwarf with a RM measurement. Its orbit normal is consistent with being aligned with the projected host star spin axis. The dotted vertical line shows the suggested dividing line between hot and cool stars by \citet{winnrm2010}. \label{fig:rmdist}} \end{figure} \subsection{Tides, Synchronization, and Kozai Emplacement}\label{sec:tides} \begin{figure} \epsscale{2.2} \plotone{tidsyn.eps} \caption{ Dimensionless combinations of physical parameters that quantify the relative time scale for orbital tidal decay (top panel) and stellar spin-orbit synchronization (bottom panel) for different binary systems, as a function of the orbital period of the system. See \S\ref{sec:tides} for an explanation and assumptions. Brown dwarfs are shown as triangles, exoplanets as squares, and low-mass stars as asterisks. KELT-1b is shown as the large star. Among known transiting exoplanets and brown dwarfs, it has the shortest characteristic time scale for orbital decay and synchronization. \label{fig:tidsyn}} \end{figure} Given the relatively large mass and short orbital period of KELT-1b, it seems probable that tides have strongly influenced the past evolution of the system, and may continue to be affecting its evolution. The literature on the influence of tides on exoplanet systems is vast (see, e.g., \citealt{rasio1996,ogilvie2004,jackson2008,leconte2010,matsumura2010,hansen2010}, for but a few examples), and there appears to be little consensus on the correct treatment of even the most basic physics of tidal dissipation, with the primary uncertainties related to where, how, and on what time scale the tidal energy is dissipated. While we are interested in evaluating the importance of tides on the evolution of the orbit of KELT-1b and the spin of KELT-1, delving into the rich but difficult subject of tides is beyond the scope of this paper. We therefore take a somewhat heuristic approach. Specifically, we construct a few dimensionless quantities that likely incorporate the primary physical properties of binary systems that determine the scale of tidal evolution, but do not depend on the uncertain physics of energy dissipation. Specifically, we define, \bea {\cal T}_a & \equiv & \frac{M_*}{M_P} \left(\frac{a}{R_*}\right)^5,\qquad {\rm and}\\ {\cal T}_{\omega_*} & \equiv & \left(\frac{M_*}{M_P}\right)^2 \left(\frac{a}{R_*}\right)^3. \label{eqn:tfacs} \eea For some classes of theories of tidal dissipation and under some assumptions, ${\cal T}_a$ is proportional to the e-folding timescale for decay of the orbit, and ${\cal T}_{\omega_*}$ is proportional to the timescale for synchronization of the spin of the star with the companion orbital period. It is worthwhile to note that for transiting planet systems the combinations of parameters $M_P/M_*$ and $a/R_*$ are generally much better determined than the individual parameters. In particular, the ratio of the mass of the planet to that of the star is closely related to the RV semi-amplitude $K$, whereas $a/R_*$ is closely related to the ratio of the duration of the transit to the period \citep{winn2010}. Figure \ref{fig:tidsyn} shows ${\cal T}_a$ and ${\cal T}_{\omega_*}$ as a function of orbital period for the sample of transiting exoplanets, brown dwarfs, and low-mass stars discussed previously. KELT-1b has shorter timescales than nearly the entire sample of systems, with the exception of a few of the low-mass stars. We therefore expect tidal effects to be quite important in this system. As a specific example, under the constant time lag model \citep{hut1981,matsumura2010}, and assuming dissipation in the star, zero eccentricity, zero stellar obliquity, and a slowly rotating star, the characteristic time scale for orbital decay due to tides is $\tau_{decay} \equiv a/|{\dot a}| = (12\pi)^{-1} Q_*' {\cal T}_a P$, where $Q_*'$ is related to the dimensionless tidal quality factor. For KELT-1b, ${\cal T}_a \sim 3 \times 10^{4}$, and so $\tau_{decay}\sim 0.3~{\rm Gyr}$ for $Q_*'=10^8$, clearly much shorter than the age of the system. Similarly, the time scale for spinning up the star by the companion is $\tau_{synch} \equiv \omega_*/|\dot \omega_*| \propto Q_*'{\cal T}_{\omega_*} P$ \citep{matsumura2010}, and so is also expected to be short compared to the age of the system. Given the expected short synchronization time scale and the fact that the expected time scale for tidal decay is shorter than the age of the system, it is interesting to ask whether or not the system has achieved full synchronization, thus ensuring the stability of KELT-1b. The measured projected rotation velocity of the star is $\vsini=56\pm 2~\kms$, which given the inferred stellar radius corresponds to a rotation period of $P_* = 2\pi R_*\sin I_*/\vsini = [1.322\pm 0.053]\sin I_*$~days, which differs from the orbital period of KELT-1b by $\sim 2\sigma$ for $I_*=90^\circ$. This is suggestive that the system is indeed synchronized. The small discrepancy could either be due to a slightly underestimated uncertainty on $\vsini$, or the host could be moderately inclined by $I_* \sim [67\pm 7]^\circ$. However, one might expect the obliquity of the star to be realigned on roughly the same time scale as the synchronization of its spin \citep{matsumura2010}. The stellar inclination can also be constrained by the precise shape of the transit light curve: lower inclinations imply higher rotation velocities, and thus increased oblateness and gravity brightening \citep{barnes2009}. Ultimately, the inclination is limited to $I_* \ga 10^\circ$ in order to avoid break up. We can also ask, given the known system parameters, if the system is theoretically expected to be able to achieve a stable synchronous state. A system is ``Darwin stable'' \citep{darwin1879,hut1980} if its total angular momentum, \be L_{tot} = L_{orb} + L_{\omega,*} + L_{\omega,P} \label{eqn:ltot} \ee is more than the critical angular momentum of \be L_{crit} \equiv 4\left[\frac{G^2}{27}\frac{M_*^3M_P^3}{M_*+M_P}(C_*+C_P)\right]^{1/4}, \label{lcrit} \ee where $L_{orb}$ is the orbital angular momentum, $L_{\omega_*}$ is the spin angular momentum of the star, $L_{\omega,P}$ is the spin angular momentum of the planet, and $C_*=\alpha_* M_* R_*^2$ and $C_P=\alpha_P M_P R_P^2$ are the moments of inertia of the star and planet, respectively \citep{matsumura2010}. Since $C_P/C_* \sim (M_P/M_*)(R_P/R_*)^2 \sim 10^{-3}$, the contribution from the planet spin to the total angular momentum is negligible. We find $L_{tot}/L_{crit}=1.029 \pm 0.014$, marginally above the critical value for stability. In addition, we find $(L_{\omega,*}+L_{\omega,P})/L_{orb} = 0.154 \pm 0.006$, which is smaller than the maximum value of $1/3$ required for a stable equilibrium \citep{hut1980}. Curiously, if we assume the star is already tidally synchronized, we instead infer $(L_{\omega,*}+L_{\omega,P})/L_{orb}=0.167 \pm 0.004$, i.e., remarkably close to exactly one-half the critical value of 1/3. Two additional pieces of information potentially provide clues to the evolutionary history of this system: the detection of a possible tertiary (\S\ref{sec:keckao}; Fig.\ref{fig:aoimage}), and the measurement of the RM effect (Fig.~\ref{fig:RM}), demonstrating that KELT-1 has small projected obliquity. If the nearby companion to KELT-1 is indeed bound, it could provide a way of emplacing KELT-1b in a small orbit via the Kozai-Lidov mechanism \citep{kozai1962,lidov1962}. If KELT-1b were originally formed much further from its host star, and on an orbit that was significantly misaligned with that of the putative tertiary, then its orbit might subsequently be driven to high eccentricity via secular perturbations from the tertiary \citep{holman1997,lithwick2011,katz2011}. If it reached sufficiently high eccentricity such that tidal effects became important at periastron, the orbit would be subsequently circularized at a relatively short period \citep{fabrycky2007,wu2007,socrates2012}. Nominally, one might expect the orbit of KELT-1b to be then left with a relatively large obliquity \citep{naoz2011}. The measured projected obliquity is $\la 16$ degrees, implying that either the current true obliquity is small, or the star is significantly inclined (i.e., $I_* \sim 0$). However, if the star is significantly inclined, then the system cannot be synchronized. Perhaps a more likely alternative is that, after emplacement by the tertiary and circularization of the orbit, the system continued to evolve under tidal forces, with KELT-1b migrating inward to its current orbit while damping the obliquity of KELT-1 and synchronizing its spin period. Clearly, detailed simulations are needed to establish whether or not this scenario has any basis in physical reality. \subsection{Comparison to Theoretical Models of Brown Dwarfs}\label{sec:models} Transiting brown dwarfs provide one of the only ways to test and calibrate models of BD structure and evolution, which are used to interpret observations of the hundreds of free floating brown dwarfs for which no direct measurement of mass and radius is possible. Given that only 5 transiting brown dwarfs with radius measurements were previously known, KELT-1b potentially provides another important test of these models. Figure \ref{fig:mr} shows the mass-radius relation for the known transiting companions to main-sequence stars with companion masses in the range $10-100~M_J$. Being close to the minimum in the brown dwarf desert, the mass of KELT-1b begins to fill in the dearth of know systems between $\sim 20-60~\mjup$. Furthermore, the formal uncertainty in its radius is only $\sim 2.5\%$, thereby allowing for a stringent test of models. In contrast, the two transiting BDs with similar masses, CoRoT-3b \citep{deleuil2008} and KOI-423b \citep{bouchy2011b}, have much larger radius uncertainties, presumably due to the relative faintness of the host stars. \begin{figure} \epsscale{1.2} \plotone{mr.eps} \caption{Radius versus mass for the known transiting companions to main-sequence stars with companion masses in the range $10-100~M_J$ that have measured radii. An estimate of the deuterium burning limit \citep{spiegel2011} is shown as the vertical dotted line, and the hydrogen burning limit is shown as the vertical dashed line. Brown dwarfs are shown as triangles, exoplanets as squares, and low-mass stars as asterisks. KELT-1b is shown as the large star. Predicted radii as a function of mass for isolated objects from the isochrones of \citet{baraffe2003} are shown for an age of 5 Gyr (dashed), 1 Gyr (dotted), and 0.5 Gyr (long dashed); the true age of the KELT-1 system is almost certainly between $1-5$ Gyr. Although stellar insolation is likely to increase the radii at fixed mass, \citet{bouchy2011a} predict that the effect is small. KELT-1b therefore has an anomalously large radius. \label{fig:mr}} \end{figure} Evolutionary models for isolated BDs generally predict that young ($\sim 0.5$ Gyr) objects in the mass range $10-100~\mjup$ should have radii of $\sim \rjup$ (see the models of \citealt{baraffe2003} in Fig.~\ref{fig:mr}). As these objects cool, however, their radii decrease, particularly for masses between 50 and 80 $\mjup$. After $\sim 1$~Gyr, all isolated objects with mass between 20-80~$\mjup$ are predicted to have radii $<\rjup$. The radius we measure for KELT-1b is $R_P = 1.110_{-0.022}^{+0.032}~\rjup$, which, at a mass of $M_P=27.23_{-0.48}^{+0.50}~\mjup$, is $\sim 7~\sigma$ and $\sim 10~\sigma$ larger than the radius predicted by \citet{baraffe2003} for ages of 1 Gyr and 5 Gyr, respectively. KELT-1b is strongly irradiated, which in principle can delay its cooling and contraction. However, \citet{bouchy2011a} predict that the effect of insolation is small for brown dwarfs in this mass range, although their models were for a much more modest insolation corresponding to an equilibrium temperature of 1800~K (versus $\sim 2400$K for KELT-1b). Therefore, given the estimated $1.5-2$ Gyr age of the system, KELT-1b is likely to be significantly inflated relative to predictions. \begin{figure} \epsscale{1.2} \plotone{rprsv.eps} \caption{ Transit depth assuming no limb darkening, i.e., $(R_P/R_*)^2$, as a function of the apparent $V$ magnitude of the host star for a sample of transiting systems. Brown dwarfs are shown as triangles, exoplanets as squares, and low-mass stars as asterisks. KELT-1b is shown as the large star. All else being equal, objects in the top left provide the best targets for follow-up. KELT-1b has a similar transit depth as the other known transiting brown dwarfs, but is significantly brighter. Also labeled are some other benchmark systems. KELT-2b ($M_P \sim 1.5~\mjup$) is shown as a large cross (Beatty et al., in preparation). \label{fig:rprsv}} \end{figure} Using the benchmark double transiting BD 2M0535$-$05, \citet{gmc2009} explore models in which brown dwarfs have large spots, which reduce the flux from their surface, thereby decreasing their effective temperatures and increasing their radii relative to those without spots (see also \citet{bouchy2011a}). They find that these can lead to significantly inflated radii, but only for large spot filling factors of $\sim 50\%$, and for relatively young ($\sim 0.5$ Gyr) systems. However, a detailed spectroscopic analysis of that system by \citet{mohanty2010} and \citet{mohanty2012}, shows that surface spots cannot be present with such a large filling factor, and thus favor global structural effects such as strong internal magnetic fields (e.g., \citealt{mullan2010}). Many other mechanisms have been invoked to explain the inflated radii of some giant exoplanets (see \citealt{fortney2010} for a review), however it is not clear which, if any, of the many mechanisms that have proposed may also be applied to inflated brown dwarfs. We would be remiss if we did not question whether we were erroneously inferring a large radius for the planet. In the past, such situations have arisen when there is a discrepancy between the constraint on the stellar density from the light curve and the constraint on the stellar surface gravity of the star from spectroscopy (e.g., \citealt{johnskrull2008,winn2008}). In our case, we find no such tension. The parameters of the star inferred from the spectroscopic data alone are in nearly perfect agreement with the results from the global analysis of the light curve, RV data, and spectroscopic constraints. We note that the effect of allowing a non-zero eccentricity also has a negligible effect on the inferred planetary radius. Finally, we reiterate that the faint companion detected in AO imaging (\S\ref{sec:keckao}), which is unresolved in our follow-up photometry, has a negligible effect on our global fit and inferred parameters. Therefore, we believe our estimate of $R_P$ is likely robust. We conclude by noting that there is a need for predictions of the radii of brown dwarfs for a range of ages and stellar insolations, and it would be worthwhile to explore whether or not the inflation mechanisms that have been invented to explain anomalously large giant planets might work for much more massive and dense objects as well. \subsection{Prospects for Follow Up} Figure \ref{fig:rprsv} compares the transit depth and apparent visual magnitude of the KELT-1 system ($\delta \sim 0.6\%$, $V=10.7$) to the sample of transiting systems collected in \S\ref{sec:comp} with available $V$ magnitudes. KELT-1 is not particularly bright compared to the bulk of the known transiting exoplanet hosts. However, it is significantly brighter than the hosts of all known transiting brown dwarfs; the next brightest is WASP-30 \citep{anderson2011}, which is $\sim 1.2$ magnitudes fainter. On the other hand, the depth of the KELT-1b transit is similar to that of the other known brown dwarfs. The prospects for follow-up of KELT-1b are exciting, not only because of the brightness of the host, but also because of the extreme nature of the system parameters, in particular the relatively short orbital period, relatively large stellar radius, and relatively large amount of stellar irradiation received by the planet. Following \citet{mazeh2010} and \citet{faigler2011}, we can estimate the amplitudes of ellipsoidal variations $A_{ellip}$, Doppler beaming $A_{beam}$ (see also \citealt{loeb2003}), reflected light eclipses and phase variations $A_{ref}$, and thermal light eclipses and phase variations $A_{therm}$, \bea A_{beam} &=& \alpha_{beam}4\left(\frac{K}{c}\right) \sim 5.7\alpha_{beam}\times 10^{-5}\\ A_{ref} &=& \alpha_{ref} \left(\frac{R_P}{a}\right)^2\sim 4.6\alpha_{ref}\times 10^{-4}\\ A_{ellip} &=& \alpha_{ellip}\frac{M_P}{M_*} \left(\frac{R_*}{a}\right)^3 \sim 4.1\alpha_{ellip}\times 10^{-4} \\ A_{therm} &=& \alpha_{therm} \left(\frac{R_P}{R_*}\right)^2 \left(\frac{R_*}{a}\right)^{1/2} \sim 3.2 \alpha_{therm}\times 10^{-3}, \eea where the expression for $A_{therm}$ assumes observations in the Rayleigh-Jeans tail of both objects, and the expression for $A_{ellip}$ assumes an edge-on orbit. The dimensionless constants $\alpha$ are defined in \citet{mazeh2010}, but to make contact with the secondary eclipse analysis in \S\ref{sec:second} we note that $\alpha_{ref}=A_g$ and $\alpha_{therm}=[f'(1-A_B)]^{1/4}$. All of these constants are expected to be of order unity, except for $\alpha_{ref}$, which may be quite low for strongly irradiated planets, depending on wavelength \citep{burrows2008}. Based on previous results, all of these effects with the possible exception of Doppler beaming are likely to be detectable with precision photometry (see, e.g., \citealt{cowan2012}). For ellipsoidal variations in particular, we expect $\alpha_{ellip} \sim 2$ and thus a relatively large amplitude of $A_{ellip} \sim 10^{-3}$. Furthermore, the detection of all these signals is facilitated by the short orbital period for KELT-1b. The prospects for transmission spectroscopy are probably poorer, given the relatively small planet/star radius ratio ($\sim 0.078$) and more importantly the large surface gravity for KELT-1b. For the optimistic case of $T_{eq}\simeq 2400$K assuming zero albedo and perfect redistribution, the scale height is only $H\sim kT/(\mu m_H g_P) \sim 16~{\rm km}$, and thus will only lead to changes in the transit depth of order $\sim 2H/R_P \sim 0.04\%$. We have presented the discovery of KELT-1b, the first transiting low-mass companion from the wide-field Kilodegree Extremely Little Telescope-North (KELT-North) transit survey. The host star KELT-1 is a mildly evolved, solar-metallicity, rapidly-rotating, mid-F star with an age of $\sim 1.5-2$ Gyr located at a distance of $\sim 260$~pc. The transiting companion is a low-mass brown dwarf or supermassive planet with mass $\sim 27~\mjup$, on a very short period, circular orbit of $P \sim 1.2$~days. In many ways, the KELT-1 system is quite unusual and extreme: KELT-1b receives a large amount of stellar insolation, is inflated relative to theoretical predictions, and raises strong tides on its host. The obliquity of KELT-1 is consistent with zero, and there is evidence that the spin of KELT-1 has been synchronized with the orbital period of KELT-1. Finally, there is a likely M-dwarf stellar companion to the KELT-1 system with a projected separation of $\sim 150$~AU. As the first definitively inflated transiting brown dwarf, KELT-1b demonstrates the need for models of brown dwarfs subject to a range of stellar insolations. A plausible formation scenario for this system posits that KELT-1b formed on a much wider orbit, and was driven to a smaller semimajor axis by the tertiary via the Kozai-Lidov mechanism. The system then continued to evolve under strong tidal forces, with KELT-1b migrating inward to its current orbit, while damping the obliquity of KELT-1 and synchronizing its spin period. The future evolution of the KELT-1 system may be spectacular. As KELT-1 continues to evolve and its radius increases, so will the tides raised on it by KELT-1b. Assuming KELT-1 is and remains tidally locked, as it cools it will develop a deep convective envelope, but be forced to rotate at an ever increasing rate. In $\sim 2$ Gyr, KELT-1 will have roughly the temperature of sun, but with a radius of $\sim 2~\rsun$ and a rotational velocity of $\sim 100~\kms$. At this point, KELT-1 will likely become an active RS CVn star \citep{walter1981}. Eventually, as KELT-1 reaches the base of the giant branch, it will swallow KELT-1b whole, likely resulting in a bright UV/X-ray and optical transient \citep{metzger2012}. | 12 | 6 | 1206.1635 |
1206 | 1206.4821_arXiv.txt | Nuclear matter and compact neutron stars are studied in the framework of the non-linear derivative (NLD) model which accounts for the momentum dependence of relativistic mean-fields. The generalized form of the energy-momentum tensor is derived which allows to consider different forms of the regulator functions in the NLD Lagrangian. The thermodynamic consistency of the NLD model is demonstrated for arbitrary choice of the regulator functions. The NLD approach describes the bulk properties of the nuclear matter and compares well with microscopic calculations and Dirac phenomenology. We further study the high density domain of the nuclear equation of state (EoS) relevant for the matter in $\beta$-equilibrium inside neutron stars. It is shown that the low density constraints imposed on the nuclear EoS and by the momentum dependence of the Schr\"odinger-equivalent optical potential lead to a maximum mass of the neutron stars around $M \simeq 2 M_{\odot}$ which accommodates the observed mass of the J1614-2230 millisecond radio pulsar. | Introduction} Relativistic hadrodynamics (RHD) of interacting nucleons and mesons provide a simple and successful tool for the theoretical description of different nuclear systems such as nuclear matter, finite nuclei, heavy-ion collisions and compact neutron stars~\cite{Serot:1984ey}. Starting from the pioneering work of Duerr~\cite{Duerr:1956zz}, simple RHD Lagrangians have been introduced~\cite{Walecka:1974qa,Serot:1997xg} and since then many different extensions of RHD approach, which rely on relativistic mean-field (RMF) approximation, have been developed. They describe the saturation mechanism in nuclear matter and generate a natural mechanism for the strong spin-orbit force in nuclei. An energy dependence of the Schr\"{o}dinger-equivalent optical potential~\cite{Cooper:1993nx,Hama:1990vr} is thereby included as a consequence of a relativistic description. However, when using the standard RHD Lagrangian in RMF approximation, the nucleon selfenergies become simple functions of density only, and do not depend explicitly on momentum of the nucleon. As a consequence a linear energy dependence of the Schr\"{o}dinger-equivalent optical potential with a divergent behavior at high energies arises~\cite{Weber:1992qc}. This well-known feature contradicts Dirac phenomenology~\cite{Cooper:1993nx,Hama:1990vr,Typel2002299}. To solve this issue one may go beyond the mean-field approximation in a quantum field theoretical framework by a systematic diagrammatic expansion of nucleon selfenergies. For instance, in Dirac-Brueckner-Hartree-Fock (DBHF)~\cite{Haar:1986ii,Brockmann:1996xy,Muther2000243} calculations the nucleon selfenergies indeed depend on both the density and single particle momentum. They reproduce the empirical saturation point of nuclear matter as well as the energy dependence of the optical potential at low energies. However, the DBHF approach has its apparent limitations at high energies and densities relevant, for instance, in heavy-ion collisions where its application within transport theory turns out to be intricate~\cite{Botermans1990115,Buss:2011mx}. Also the thermodynamic consistency of the DBHF calculations is not obvious~\cite{Hugenholtz:1958}. As an alternative approach to {\it ab-initio} DBHF calculations for the nuclear many-body systems a phenomenological treatment of the problem in the spirit of the RMF approximation is still considered as a powerful tool. However, the simple Lagrangian of RHD~\cite{Walecka:1974qa,Serot:1997xg} has to be further modified for a quantitative description of static nuclear systems such as nuclear matter and/or finite nuclei. Therefore, it is mandatory to introduce new terms, {\it e.g.}, including non-linear self interactions of the scalar~\cite{Boguta:1977xi} and/or vector~\cite{Sugahara1994557} meson fields, or to modify existing contributions in the Lagrangian, {\it e.g.}, introducing density dependent meson-nucleon couplings~\cite{PhysRevLett.68.3408,Fuchs:1995as,Typel:1999yq}. The model parameters have to be then fitted to properties of nuclear matter and/or atomic nuclei, since, they cannot be derived in a simple manner from a microscopic description. The momentum dependence of in-medium interactions becomes particularly important in description of nuclear collision dynamics such as heavy-ion collisions. Indeed, analyses of proton-nucleus scattering data \cite{Cooper:1993nx,Hama:1990vr} show that the proton-nucleus optical potential starts to level off already at incident energies of about $300$ MeV. Thus, other RMF approaches have been developed by including additional non-local contributions, {\it i.e.}, by introducing Fock-terms, on the level of the RMF selfenergies leading to a density and momentum dependent interactions~\cite{Weber:1992qc}. However, such a treatment is not covariant and also its numerical realization in actual transport calculations is rather difficult~\cite{Weber:1992qc}. Another approach has been proposed in~\cite{Zimanyi:1990np} and more recently in~\cite{Typel:2002ck,Typel:2005ba} by introducing higher order derivative couplings in the Lagrangian of RHD. In Ref.~\cite{Zimanyi:1990np} such gradient terms have been studied with the conclusion of a softening of the nuclear EoS. In another study of Ref.~\cite{Typel:2005ba} both the density dependence of the nuclear EoS and the energy dependence of the optical potential have been investigated. While the modified interactions of meson fields with nucleons explain the empirical energy dependence of the optical potential, a stiff EoS at high densities results from an introduction of an explicit density dependence of the nucleon-meson couplings with additional parameters. The impact of momentum dependent RMF models on nuclear matter bulk properties and particularly on the high density domain of EoS relevant for neutron stars is presently less understood. The purpose of the present work is to develop a relativistic and thermodynamically consistent RMF model, which provides the correct momentum dependence of the nucleon selfenergies and agrees well with available empirical information on nuclear matter ground state, in a self consistent Lagrangian framework. Some steps in this direction have been already done in Refs.~\cite{Gaitanos:2011yb,Gaitanos:2011ej,Gaitanos:2009nt} where the concept of non-linear derivative meson-nucleon Lagrangian has been introduced. However, the calculations of Refs.~\cite{Gaitanos:2011yb,Gaitanos:2011ej,Gaitanos:2009nt} were based on a particular exponential form of the regulators in the RHD Lagrangian and a detailed study of nuclear matter ground state properties has not been done. In the present work the generalized form of the energy-momentum tensor in the NLD model is derived and allows to consider different regulator functions in the Lagrangian. The thermodynamic consistency of the NLD model is demonstrated for arbitrary choice of the regulators. A thorough study of the properties of nuclear matter around saturation density is further performed. The model describes the bulk properties of the nuclear matter and compares well with microscopic calculations and Dirac phenomenology. We also investigate the high density region of the NLD EoS relevant for the neutron stars. It is found that the low density constraints imposed on the nuclear matter EoS and by the momentum dependence of the Schr\"odinger-equivalent optical potential lead to a maximum mass of the neutron stars around $M \simeq 2 M_{\odot}$. It is demonstrated that the high density pressure-density diagram as extracted from astrophysical measurements~\cite{Ozel:2010fw,Steiner:2010fz} can be well described with nucleonic degrees of freedom only. | 12 | 6 | 1206.4821 |
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1206 | 1206.1773_arXiv.txt | { The Cosmological Microwave Background (CMB) is of premier importance for the cosmologists to study the birth of our universe. Unfortunately, most CMB experiments such as COBE, WMAP or Planck do not provide a direct measure of the cosmological signal; CMB is mixed up with galactic foregrounds and point sources. For the sake of scientific exploitation, measuring the CMB requires extracting several different astrophysical components (CMB, Sunyaev-Zel'dovich clusters, galactic dust) form multi-wavelength observations. Mathematically speaking, the problem of disentangling the CMB map from the galactic foregrounds amounts to a component or source separation problem. In the field of CMB studies, a very large range of source separation methods have been applied which all differ from each other in the way they model the data and the criteria they rely on to separate components. Two main difficulties are i) the instrument's beam varies across frequencies and ii) the emission laws of most astrophysical components vary across pixels. This paper aims at introducing a very accurate modeling of CMB data, based on sparsity, accounting for beams variability across frequencies as well as spatial variations of the components' spectral characteristics. Based on this new sparse modeling of the data, a sparsity-based component separation method coined Local-Generalized Morphological Component Analysis (L-GMCA) is described. Extensive numerical experiments have been carried out with simulated Planck data. These experiments show the high efficiency of the proposed component separation methods to estimate a clean CMB map with a very low foreground contamination, which makes L-GMCA of prime interest for CMB studies.} | \label{sec:conc} The estimation of a high precision CMB map featuring low noise and low foreground contamination is of crucial interest for the astrophysical community. This problem is customarily tackled in the framework of component separation. As any estimation problem, the accurate modeling of the data is essential. However, a close look at the astrophysical phenonema at play in the CMB data such as WMAP or Planck reveals that the linear mixture model used so far by common component separation methods in cosmology does not hold. First, the variation of the beam across scale is seldom accounted for which highly limits the performances of these component separation methods especially at small scales. More importantly, the linear mixture model does not afford enough degrees of freedom to precisely capture the complexity of the data including the variability across space of the emission law of some of the components of interest. To alleviate these limitations, this paper introduces a new modeling of the components' mixtures using a multiscale and local decomposition of the data in the wavelet domain. We introduced a novel sparsity-based coined L-GMCA (Local - Generalized Morphological Component Analysis) which makes profit of the proposed local/multiscale mixture model. In the proposed framework, wavelet-based multiscale analysis allows for a precise modeling of the beam evolution across channels. Capturing the variations across pixels of the emissivity of components is carried out by partitioning each wavelet scale with adaptive patch sizes. Extensive numerical experiments have been carried out which show the superiority of the proposed modeling and separation technique to provide a clean, low-foreground CMB map estimate. More precisely, we showed that the local and multiscale modeling allows for improved separation results whether it is used with separation techniques as different as ILC and GMCA. Additionally, the numerical experiments enlightens the dramatic positive impact of the use of sparsity in L-GMCA to provide less galactic foreground contamination as well as significantly lower non-gaussianity levels. | 12 | 6 | 1206.1773 |
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1206 | 1206.3540_arXiv.txt | *{The Earth is continuously showered by charged cosmic ray particles, naturally produced atomic nuclei moving with velocity close to the speed of light. Among these are ultra high energy cosmic ray particles with energy exceeding $5\times 10^{19}$ eV, which is ten million times more energetic than the most energetic particles produced at the Large Hadron Collider. Astrophysical questions include: what phenomenon accelerates particles to such high energies, and what sort of nuclei are energized? Also, the magnetic deflection of the trajectories of the cosmic rays makes them potential probes of galactic and intergalactic magnetic fields. We develop a Bayesian hierarchical model that can be used to compare different association models between the cosmic rays and source population, using Bayes factors. A measurement model with directional uncertainties and accounting for non-uniform sky exposure is incoporated into the model. The methodology allows us to learn about astrophysical parameters, such as those governing the source luminosity function and the cosmic magnetic field.} | \label{sec:1} Since the Pierre Auger Observatory (PAO) initiated observations in 2004 it has detected 14 ultra high energy cosmic rays (UHECRs) with energy $\geq 55$ Eev in period 1--January 1, 2004 - May 26, 2006, 13 UHECRs in period 2-- May 27, 2006 - August 31, 2007, and 42 UHECRs in period 3-- September 1, 2007 - December 31, 2009. The energy threshold of 55 Eev was chosen by using period 1 data \cite{PAO}. These CR particles interact with the cosmic microwave background, and according to GZK limit, CRs with energy $\simgreat$ 60 Eev should come from sources within 200 Mpc \cite{PAO}. We consider the 17 active galactic nuclei (AGNs) in the volume-complete (to 15 Mpc) catalog of \cite{AGN} as candidate sources. We use a Bayesian hierarchical model to compare 3 models, M$_0$: only isotropic background source, M$_1$: isotropic background + 17 AGNs, M$_2$: isotropic background + 2 AGNs (Centaurus A and NGC 4945--the two closest AGNs) for the UHECRs from the three periods. | 12 | 6 | 1206.3540 |
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1206 | 1206.7069_arXiv.txt | The Haumea family is currently the only identified collisional family in the Kuiper belt. We numerically simulate the long-term dynamical evolution of the family to estimate a lower limit of the family's age and to assess how the population of the family and its dynamical clustering are preserved over Gyr timescales. We find that the family is not younger than 100 Myr, and its age is at least 1 Gyr with 95\% confidence. We find that for initial velocity dispersions of $50-400$ ms$^{-1}$, approximately $20-45$\% of the family members are lost to close encounters with Neptune after 3.5 Gyr of orbital evolution. We apply these loss rates to two proposed models for the formation of the Haumea family, a graze-and-merge type collision between two similarly sized, differentiated KBOs or the collisional disruption of a satellite orbiting Haumea. For the graze-and-merge collision model, we calculate that $>85\%$ of the expected mass in surviving family members within $150$~ms$^{-1}$ of the collision has been identified, but that one to two times the mass of the known family members remains to be identified at larger velocities. For the satellite-break-up model, we estimate that the currently identified family members account for $\sim50\%$ of the expected mass of the family. Taking observational incompleteness into account, the observed number of Haumea family members is consistent with either formation scenario at the $1 \sigma$ level, however both models predict more objects at larger relative velocities ($>150$~ms$^{-1}$) than have been identified. | \label{s:Introduction} The Haumea (2003 EL$_{61}$) collisional family was discovered by \citet{Brown2007} who noted that Haumea and five other Kuiper belt objects (KBOs) shared a spectral feature that is indicative of nearly pure water ice on the surfaces of the bodies. These six KBOs, along with four additional family members identified by \citet{Schaller2008}, \citet{Snodgrass2010}, and \citet{Ragozzine2007}, can all be dynamically linked to Haumea, and there do not appear to be any dynamically unrelated KBOs that share this spectral feature. Aside from being spectrally linked to these other KBOs, Haumea itself shows signs of its collisional past. Despite having a nearly pure water ice surface, Haumea's density is $\sim2.6$ g cm$^{-3}$ \citep{Rabinowitz2006}, which is higher than expected for typical assumed ice/rock ratios in the Kuiper belt \citep{Brown2008}; one way to achieve this higher density is to have a catastrophic collision between a differentiated proto-Haumea and another KBO in which proto-Haumea loses a substantial fraction of its water ice mantle \citep{Brown2007}. This scenario is supported by the presence of at least two water ice satellites \citep{Barkume2006,Ragozzine2009}. Haumea also has an elongated shape and a very short spin period of $\sim4$ hours that is unlikely to be primordial \citep{Rabinowitz2006,Lacerda2007}. \citet{Ragozzine2007} examined the dynamical connections between the identified Haumea family members. These connections are made by first estimating the orbit of the center of mass of the colliding bodies, and then estimating the ejection velocities of each family member relative to the collision's center of mass. The ejection velocity is given by \begin{equation}\label{eq:dv} \Delta \vec{v}= \vec{v} - \vec{v}_{cm} \end{equation} where $\vec{v}_{cm}$ is the estimated collision's center-of-mass velocity. Because Haumea is by far the largest remnant from the collision, its orbit immediately after the collision should have nearly coincided with the center-of-mass orbit. However, Haumea is currently located at the boundary of the 12:7 mean motion resonance (MMR) with Neptune; over long timescales, the chaotic zone of this resonance causes a random walk of the proper elements such that Haumea's current orbit may be significantly distant from its post-collision orbit. \citet{Ragozzine2007} estimate the center-of-mass collision orbit by minimizing the sum of the relative speeds of all family members, assuming that Haumea's semimajor axis and its Tisserand parameter with respect to Neptune are both conserved during its chaotic evolution; they then use Haumea's present distance from the collision's center-of-mass orbit, together with a calculation of its chaotic diffusion rate, to estimate the age of the collisional family to be $3.5\pm2$ Gy. Given the exceedingly low collision probabilities for objects large enough to form the Haumea family in the current Kuiper belt, the family is likely to be old. However, the family probably cannot have formed in the primordial, much more massive Kuiper belt, because whatever caused the mass of the Kuiper belt to be depleted (by an estimated 2 or 3 orders of magnitude) would have also destroyed the dynamical coherence of the family~\citep{Levison2008}. The high inclination ($\sim27^{\circ}$) of the family also argues against a primordial origin, because such large inclinations are probably products of the excitation and mass depletion of the Kuiper belt. Thus, it appears that the Haumea family-forming collision occurred near the end of the primordial, high-mass phase of the Kuiper belt. Several of the largest KBOs show evidence of their collisional past (see review by \citet{Brown2008}), but the Haumea family is the only collisional family that has been identified in the Kuiper belt. The dynamical connections between the members of the family allow us to place some constraints on the type of collision that formed the family and also constrain the age of the family as being old, but probably not primordial. These characteristics make the Haumea family an excellent probe of the collisional environment in the Kuiper belt following the excitation and mass depletion event; understanding the type of collision that created the family (especially the relative sizes and speeds of the impactor and target) would provide valuable insight into the size and orbital distribution of the Kuiper belt at the time of the collision (see discussions of this in \citet{Marcus2011} and \citet{Levison2008}). Proposed models for the formation of the Haumea family have attempted to reproduce the family's relatively small velocity dispersion ($\sim150$~ms$^{-1}$) and to explain the compositional and orbital characteristics of the family. However, the orbits of the family members have been sculpted by several gigayears of dynamical evolution. In this paper we use numerical simulations to determine how this orbital evolution affects the dynamical coherence of the family. In Section~\ref{s:sims}, we determine the loss rates for the family, which depend on the initial velocity dispersion from the collision, and we determine how the velocity dispersion of the surviving family members is altered over time; from these simulations, we also obtain a hard lower limit for the age of the family. In Section~\ref{s:formation_models}, we apply these results to the family-formation models of \citet{Leinhardt2010} (a graze-and-merge type collision between two similarly sized, differentiated KBOs) and \citet{Schlichting2009} (the collisional disruption of a satellite orbiting Haumea), and we compare the predictions from these two formation models to the current observations of the family. Section~\ref{s:conclusions} provides a summary of our results and conclusions. | \label{s:conclusions} After accounting for 3.5 Gyr of dynamical evolution and accounting for the observational incompleteness of the known family, both of the proposed Haumea family formation models we have examined here \citep{Leinhardt2010,Schlichting2009} are consistent with the total number of observed family members. There is, however, a significant difference between the observed values of $\Delta v$ and those predicted in the formation models; almost none ($\ll1\%$) of the synthetic families we generated for either formation scenario account for the observations of family members that all fall within $150$~ms$^{-1}$ of the collision center. The only way we find to make the velocity distributions for the formation scenarios consistent with the observations is to allow that the actual ejection velocities of some of the known family members are a factor of $\sim2$ larger than the calculated minimum values. This adjustment of the observed $\Delta v$ falls within the uncertainties for their calculation of the collision center-of-mass orbit (see our discussion in Section~\ref{ss:cmorbit}). But even with this adjustment, only $10-20\%$ of the synthetic families reproduce the observed family. This is a reasonable level of agreement given the uncertainties in the collision models and the small number of observed family members, but it is still interesting that so few family members have been identified at large $\Delta v$. In section~\ref{ss:observations} we explored how the assumptions we made about the albedos of the family members and possible correlations between fragment mass and ejection velocity could change the predictions of the graze-and-merge collision model. We find that it is difficult to alter these assumptions enough to obtain a better than 10--15\% agreement with the observations. For either formation scenario, we find that there should be an additional $\sim0.01-0.03$ $M_H$ of family members that have yet to be identified, i.e., at least as many as those already detected. If, as we discover these additional family members, the distribution of $\Delta v$ remains too heavily weighted toward the low-end, the formation models will have to be reconsidered. Recent spectroscopic and photometric surveys of the Kuiper belt designed to detect water ice have not identified any large $\Delta v$, ice-rich Haumea family members. \citet{Fraser2012} performed a photometric study of 120 objects from the major dynamical classes of the Kuiper belt using the Hubble Space Telescope (HST) and did not find any additional Haumea family members. \citet{Benecchi2011} also performed HST photometry of a large sample of Kuiper belt objects and failed to identify any new ice-rich family members. Ground based studies searching for water ice in the Kuiper belt have also failed to detect higher $\Delta v$ family members; \citet{Brown2012} detected no additional family members, and \citet{Trujillo2011} found only one new member, located near the dynamical core of the family. The extent of these surveys suggests that if there were ice-rich Haumea family members spread at large $\Delta v$ throughout the Kuiper belt, we likely should have already identified some of them. This is consistent with our comparisons of the proposed collisional models to the observed Haumea family; figure~\ref{f:simulated_detections} shows the distribution of simulated detections for a subset of our dynamically evolved graze-and-merge collisional families (see Section~\ref{ss:observations}) compared to the known Haumea family. The simulated detections span a larger parameter range of the Kuiper belt (due to the presence of higher $\Delta v$ family members) than the actual detections; \citet{Lykawka2012} also found that simulated Haumea families with $\Delta v$ consistent with existing collisional formation models would occupy a larger portion of the Kuiper belt than currently observed. In contrast to the expected $\Delta v$ distributions from the collisional models, we find that, accounting for the time evolution of the a-e-i and $\Delta v$ distributions in our simulations, all of the known family members are consistent with initial $\Delta v < 100$ ms$^{-1}$. One possible way to modify the collision models was suggested by \citet{Cook2011}, who argue that perhaps the proto-Haumea was only partially differentiated. This would allow for some of the collisional fragments to be rocky rather than primarily water ice, which would mean that they might not show the water ice spectral feature that has thus far been used as the only secure way to identify a family member; perhaps the dynamically nearby KBO 2008 AP$_{129}$, which has a $\Delta v$ of only 140~ms$^{-1}$ but shows less water ice absorption than the accepted family members (see Section~\ref{ss:observations}) represents a new class of rockier family members. A rockier composition could also make the fragments darker and therefore more difficult to detect at all. It has been similarly suggested that surface inhomogeneities on the proto-Haumea could result in collisional fragments with different compositional characteristics from those of the known family members \citep{Schaller2008}. It is unclear why, in either of these situations, there would be a preference for the less icy fragments to be dispersed at large $\Delta v$, but if the composition of the target and impactor are substantially different from those assumed in the collision simulations, that could affect the entire $\Delta v$ distribution. Another possibility is that the formation models have not adequately accounted for collisional evolution amongst the ejected fragments themselves, and that this could alter the family's size and/or velocity distribution in a significant way. Collision simulations like the ones in \citet{Leinhardt2010} are computationally very expensive and they follow the family's evolution for only a few thousand spin periods of the primary (a few hundred days in total), so the model's final size and velocity distributions are not fully evolved. The presence or absence of higher $\Delta v$ family members with future additions to the set of observed Haumea family members will determine if any of these modified scenarios should be considered, or if there is another collisional model that could better explain the family. For either the \citet{Leinhardt2010} or the \citet{Schlichting2009} models to be consistent with observations, we should find several higher $\Delta v$ family members among any new identifications. In summary, our study of the long-term dynamical evolution of the Haumea family leads to the followings conclusions. \begin{enumerate} \item The Haumea family is at least 100 Myr old. This estimate is based on the timescale to randomize the nodal longitudes of the orbital planes of the family members, as well the timescale for chaotic evolution of Haumea's eccentricity in the 12:7 MMR with Neptune. From the chaotic diffusion of Haumea's eccentricity, we can conclude with $95\%$ confidence that the family is older than 1 Gyr. \item For initial ejection velocities, $\Delta v$, in the range $50-400$~ms$^{-1}$, $20-45$\% of original Haumea family members are lost due to close encounters with Neptune over 3.5 Gyr. Most of this loss occurs at the inner edge of the family (interior to $\sim 41$ AU) and near the locations of MMRs with Neptune. A few percent of the surviving Haumea family members are expected to be found in MMRs with Neptune. The 3:2 and 7:4 MMRs are the most likely of the resonances to contain surviving members. \item Within the population of surviving and potentially recognizable family members, chaotic diffusion in orbital elements over 3.5 Gyr introduces a $50-100$~ms$^{-1}$ spread in the apparent velocities of the family relative to the collision center-of-mass orbit, with the average $\Delta v$ increasing slightly over time. \item Accounting for long-term dynamical evolution to the graze-and-merge collision model of \citet{Leinhardt2010}, we find that the currently observed family represents $>85\%$ of the expected family mass within $150$~ms$^{-1}$ of the collision center, but an additional $0.035\pm0.01$ $M_H$ (about twice the mass of the known family) remains to be identified at larger $\Delta v$. Accounting for observational incompleteness, the \citet{Leinhardt2010} model is consistent with the observations at the $\sim10\%$ confidence level. \item For the satellite breakup model of \citet{Schlichting2009}, we find that the currently observed family accounts for $\sim 50\%$ of the expected mass of the family. Most of the remaining mass should be found at $\Delta v > 150$~ms$^{-1}$. Accounting for observational incompleteness, the satellite breakup model is consistent with the observations at the $\sim20\%$ confidence level. \item Both formation models predict more family members at large $\Delta v$ than are currently observed (even allowing for a factor of $\sim2$ higher values of $\Delta v$ for the known family members due to the uncertainty in estimates of the collision center-of-mass orbit). If additional Haumea family members are identified and continue to have low $\Delta v$ ($\lesssim 200$~ms$^{-1}$), new formation models (or modifications to the existing models) will have to be considered. \end{enumerate} | 12 | 6 | 1206.7069 |
1206 | 1206.5546_arXiv.txt | A Chern-Simons coupling of a new scalar field to electromagnetism may give rise to cosmological birefringence, a rotation of the linear polarization of electromagnetic waves as they propagate over cosmological distances. Prior work has sought this rotation, assuming the rotation angle to be uniform across the sky, by looking for the parity-violating TB and EB correlations a uniform rotation produces in the CMB temperature/polarization. However, if the scalar field that gives rise to cosmological birefringence has spatial fluctuations, then the rotation angle may vary across the sky. Here we search for direction-dependent cosmological birefringence in the WMAP-7 data. We report the first CMB constraint on the rotation-angle power spectrum $C_L^{\alpha\alpha}$ for multipoles between $L=0$ and $L=512$. We also obtain a $68\%$ confidence-level upper limit of $\sqrt{C_2^{\alpha\alpha}/(4\pi)}\lesssim 1^{\circ}$ on the quadrupole of a scale-invariant rotation-angle power spectrum. | \label{sec:intro} In this work, we use the cosmic microwave background (CMB) temperature and polarization maps of the Wilkinson Microwave Anisotropy Probe (WMAP) 7-year data release \cite{Jarosik:2010iu} to search for direction-dependent cosmological birefringence (CB). CB is a postulated rotation of the linear polarization of photons that propagate through cosmological distances \cite{Carroll:1989vb}. It is present, for example, in models where a Nambu-Goldstone boson plays the role of quintessence \cite{Carroll:1998zi}, but also in models with new scalar degrees of freedom that have nothing to do with quintessence \cite{Li:2008tma,Pospelov:2008gg,Caldwell:2011pu,Finelli:2008jv}. The rotation of the polarization is a consequence of the coupling of a scalar field to the electromagnetic Chern-Simons term, such that the rotation angle $\alpha$ is proportional to the total change $\Delta\phi$ of the field $\phi$ along the photon's path. Prior to this work, a rotation angle $\alpha$ that is uniform across the sky had been sought in the CMB \cite{earlysearches}, where it would induce parity-violating TB and EB temperature/polarization correlations \cite{Lue:1998mq}. CB has also been sought in quasar data \cite{quasarsearches,Carroll:1989vb}. The tightest constraint currently comes from a combined analysis of the WMAP, Bicep \cite{Chiang:2009xsa}, and QUAD experiments \cite{Wu:2008qb}; it is $- {1.4^ \circ } < \alpha < {0.9^ \circ }$ at the $95\%$ confidence level \cite{Komatsu:2010fb}. There are, however, a number of reasons to expand the search and look for a CB angle $\alpha(\hatn)$ that varies as a function of position $\hatn$ on the sky. To begin with, a dynamical field $\phi$ that drives the rotation can have fluctuations, in which case the rotation angle varies across the sky \cite{Li:2008tma,Pospelov:2008gg,Caldwell:2011pu}. Furthermore, if $\phi$ is some massless scalar, not necessarily quintessence, its background value does not necessarily evolve, and the uniform component of the rotation angle may vanish. The only way to look for CB in this scenario is through its direction dependence. Additionally, if $\alpha(\hatn)$ is measured with high significance, the exact shape of its power spectrum provides a window into the detailed physics of the new cosmic scalar $\phi$. Currently, the strongest limit on a direction-dependent CB angle comes from AGN \cite{Kamionkowski:2010ss}, which constrain the root-variance of the rotation angle to be $\lesssim 3.7^\circ$. In previous studies \cite{Kamionkowski:2008fp, Gluscevic:2009mm}, a formalism was developed to search for anisotropic CB rotation with the CMB. The sensitivity of WMAP data to this anisotropic rotation is expected to be competitive with that from AGN \cite{Kamionkowski:2008fp,Gluscevic:2009mm,Yadav:2009eb}. However, the CMB also allows individual multipoles $C_L^{\alpha\alpha}$ to be probed---the AGN data only constrain the variance---and is sensitive to higher $L$ than AGN. The CMB also probes CB to a larger lookback time than AGN. Here we apply the formalism developed earlier to the WMAP 7-year data. Within experimental precision, we report a non-detection of a direction-dependent cosmological birefringence. We obtain an upper limit on all the rotation-angle power-spectrum multipoles $C_L^{\alpha\alpha}$ up to $L=512$. This result implies a $68\%$ confidence-level upper limit on the quadrupole of a scale-invariant power spectrum of $\sqrt{C_2^{\alpha\alpha}/(4\pi)}\lesssim 1^{\circ}$,\footnote{Here, the power spectrum is defined in the usual way, $C_L^{\alpha\alpha}\equiv \sum\limits_{M} \alpha_{LM}\alpha_{LM}^*/(2L+1)$, where a spherical-harmonic decomposition of the rotation field provides the rotation-angle multipoles, $\alpha_{LM}\equiv\int Y_{LM}^*(\hatn) \alpha(\hatn)d\hatn$.}. As a check, we also find a constraint on the uniform rotation that agrees with the results of Ref.~\cite{Komatsu:2010fb}. The rest of this paper is organized as follows. In \S\ref{sec:physics}, we review the physical mechanism for CB. In \S\ref{sec:formalism}, we revisit the full-sky formalism to search for direction-dependent rotation, and discuss its implementation. WMAP data selection, our simulations, and the tests of the analysis method are described in \S\ref{sec:data_simulations}. Results are reported in \S\ref{sec:constraints}, and we conclude in \S\ref{sec:discussion}. Appendix \ref{ax:small_alpha} contains a detailed explanation of the procedure we used to obtain an upper limit of the root-mean-squared rotation angle from the measurement of the $TE$ correlation in the data; Appendix \ref{ax:kernels} contains a discussion of the geometrical properties of the rotation-angle estimator; Appendix \ref{ax:Lfsky} details the calculation of the $L$-dependence of the fractional correction for a scale-invariant power spectrum recovered from cut-sky maps; and Appendix \ref{ax:masks} displays the analysis masks we used in this work. | \label{sec:discussion} In this work, we implement the minimum-variance quadratic-estimator formalism of Ref.~\cite{Gluscevic:2009mm} to search for direction-dependent cosmological birefringence with WMAP 7-year data. We derive the first CMB measurement of the rotation-angle power spectrum in the range $L=0-512$, finding consistency with zero at each multipole, within $3\sigma$. We estimate an upper limit on each power-spectrum multipole by simulating a suite of Gaussian sky realizations with no rotation, including symmetric beams and noise realizations appropriate for each WMAP frequency band, and also $Q$-$U$ correlations and sky cuts. We investigate the impact of foregrounds and polarized diffuse point sources on the reported constraints and come to the conclusion that they are not significant sources of systematic error for the rotation-angle estimates. Finally, we use the null result to get a $68\%$ confidence-level upper limit of $\sqrt{C_2^{\alpha\alpha}/(4\pi)}\lesssim 1^{\circ}$ on the quadrupole of a scale-independent rotation-angle power spectrum. Even though the CMB constraint turns out to be comparable to that derived from quasar measurements, the CMB analysis has a significant advantage: it provides a measurement of the rotation-angle power at each individual multipole $L$ and has better sensitivity to models with significant power at high multipoles. The same formalism we use in this work can be applied to forecast the sensitivity of upcoming and future-generation CMB satellites to detecting direction-dependent cosmological birefringence. With 7 years worth of integration time with WMAP, we are able to constrain the uniform component of the rotation to less than about a degree; it will be interesting to see the results of this analysis method applied to the upcoming data from Planck satellite \cite{Planck}, where the sensitivity to rotation angles on the order of a few arcminutes \cite{Gluscevic:2009mm} is expected. | 12 | 6 | 1206.5546 |
1206 | 1206.1139.txt | We investigate the effect of a variation of fundamental constants on primordial element production in big bang nucleosynthesis (BBN). We focus on the effect of a possible change in the nucleon-nucleon interaction on nuclear reaction rates involving the $A=5$ ($^5$Li and $^5$He) and $A=8$ ($^8$Be) unstable nuclei and complement earlier work on its effect on the binding energy of deuterium. The reaction rates for \hdp\ and \tdn\ are dominated by the properties of broad analog resonances in $^5$He and $^5$Li compound nuclei respectively. While the triple alpha process \aaag\ is normally not effective in BBN, its rate is very sensitive to the position of the ``Hoyle state" and could in principle be drastically affected if $^8$Be were stable during BBN. The nuclear properties (resonance energies in $^5$He and $^5$Li nuclei, and the binding energies of $^8$Be and D) are all computed in a consistent way using a microscopic cluster model. The n(p,$\gamma$)d, \hdp, \tdn\ and \aaag, reaction rates are subsequently calculated as a function of the nucleon-nucleon interaction that can be related to the fundamental constants. We found that the effect of the variation of constants on the \hdp, \tdn\ and \aaag\ reaction rates is not sufficient to induce a significant effect on BBN, even if $^8$Be was stable. In particular, no significant production of carbon by the triple alpha reaction is found when compared to standard BBN. We also update our previous analysis on the effect of a variation of constants on the n(p,$\gamma$)d reaction rate. | Constraints on the possible variation of fundamental constants are an efficient method of testing the equivalence principle~\cite{jpurmp,jpurevues}, which underpins metric theories of gravity and in particular general relativity. These constraints are derived from a wide variety of physical systems and span a large spectrum of redshifts and physical conditions, from the comparison of atomic clocks in the laboratory, the Oklo phenomena, to quasar absorption spectra up to a typical redshift of order $z\sim 2$ and big bang nucleosynthesis (BBN) at a redshift of order $z \sim10^{9}$. Primordial nucleosynthesis is considered a major pillar of the standard cosmological model (see e.g., Refs.~\cite{books}). Using inputs from WMAP for the baryon density~\cite{wmap7}, BBN yields excellent agreement between the theoretical predictions and astrophysical determinations for the abundances of D and \he4~\cite{bbn2,Coc04,Ioc07,cfo5,Coc12} despite the discrepancy between the theoretical prediction of \li7 and its determined abundance in halo stars \cite{cfo5}. Indeed, BBN has been used extensively to constrain deviations from the standard model framework, and in particular from general relativity, see e.g., Ref.~\cite{bbnrg}. The effects of the variation of fundamental constants on BBN predictions is difficult to model because of the intricate structure of QCD and its role in low energy nuclear reactions and because one cannot restrict the analysis to a single constant. One can, however, proceed in a two step approach: first by determining the dependencies of the light element abundances on the BBN parameters and then by relating those parameters to the fundamental constants (see Section~3.8 of Ref.~\cite{jpurmp} for an up-to-date overview). While early works have mostly focused on a single parameter such as the fine structure constant~\cite{alphaseul}, the Higgs vacuum expectation value (vev)~\cite{dixit,vvar,scherrer} or the QCD scale~\cite{lqcdseul}, in many theories which allow for the variation of fundamental parameters, often, the variation of several parameters are correlated in a model dependent way \cite{wett,co}. The variation of a fundamental parameter such as the fine structure constant will affect the BBN analysis through the proton-to-neutron mass difference and the neutron lifetime~\cite{alphaseul} as well as the deuterium binding energy~\cite{flam1} and the binding energies of other light nuclei such as tritium, helium-3 and 4, lithium-6 and 7, and beryllium-7~\cite{other-B}. These effects can in principle be used to probe the coupled variations of several parameters~\cite{bbnmultipara,Dent-et-al,Coc07}. Following our previous work~\cite{Coc07,Eks10,Luo11}, we allow for a variation of all fundamental constants and in order to reduce the arbitrariness, we focus on scenarios in which the variations of the different constants are correlated. In effectively all unification models of non-gravitational interactions, a variation in the fine structure constant is associated with the variation in other gauge couplings~\cite{wett,co}. Any variation of the strong gauge coupling $\alpha_s$ will induce a variation in the QCD scale, $\Lambda_{\rm QCD}$, as can be seen from the expression \begin{equation} \Lambda_{\rm QCD}=\mu\left(\frac{m_{\rm c}m_{\rm b}m_{\rm t}}{\mu^3}\right)^{\frac{2}{27}}\exp\left[- \frac{2\pi}{9\alpha_s(\mu)}\right] \end{equation} valid for a renormalization scale $\mu>m_{\rm t}$. In this expression $m_{\rm c,b,t}$ are the masses of the charm, bottom and top quarks. Since the masses of the quarks are proportional to the product, $hv$, of a Yukawa coupling $h$ and Higgs vacuum expectation value $v$, any variation of the Yukawa couplings will also induce a variation of $\Lambda_{\rm QCD}$. These variations can be related by \begin{equation}\label{DeltaLambda} \frac{\Delta \Lambda}{\Lambda} = R \, \frac{\Delta \alpha}{\alpha} + \frac{2}{27} \left(3 \, \frac{\Delta v}{v} + \frac{\Delta h_{\rm c}}{h_{\rm c}} + \frac{\Delta h_{\rm b}}{h_{\rm b}} + \frac{\Delta h_{\rm t}}{h_{\rm t}} \right) \,. \end{equation} The coefficient $R$ is determined by the particular grand unified theory and particle content of the theory which control both the value of $\alpha(M_{\rm GUT}) = \alpha_s(M_{\rm GUT})$ and the low energy relation between $\alpha$ and $\alpha_s$, leading to a considerable model dependence in its value~\cite{Dent,dine}. Here we shall assume a typical value, $R \sim 36$~\cite{co,Langacker}. Furthermore, in theories in which the electroweak scale is derived by dimensional transmutation, changes in the Yukawa couplings (particularly the top Yukawa, $h_t$) lead to exponentially large changes in the Higgs vev. In such cases, the Higgs expectation value is related to the Planck mass, $M_{\rm P}$, by \begin{equation} v\sim M_{\rm P}\exp\left(-\frac{2\pi c}{\alpha_{\rm t}}\right) \end{equation} where $c$ is a constant of order unity, and $\alpha_{\rm t}=h_{\rm t}^2/4\pi$. Thus we can write, \begin{equation}\label{Deltav} \frac{\Delta v}{v} \equiv S\, \frac{\Delta h}{h} \,, \end{equation} and, as in Ref.~\cite{Coc07}, we take $S\sim240$, though there is considerable model-dependence in this value as well. For example, in supersymmetric models, $S$ can be related to the sensitivity of the Z gauge boson mass to the top Yukawa, and may take values anywhere from about 80 to 500~\cite{santoso}. This dependence gets translated into a variation in all low energy particle masses~\cite{dixit}. In addition, in many string theories, all gauge and Yukawa couplings are determined by the expectation value of a dilaton. Therefore, once we allow $\alpha$ to vary, virtually all masses and couplings are expected to vary as well, typically much more strongly than the variation induced by the Coulomb interaction alone. The use of coupled variations has led to significantly improved constraints in a wide range of environments ranging from big bang nucleosynthesis~\cite{co,bbnmultipara,Coc07,flam1,other-B,fl1,Ber10,landau,grant}, the Oklo reactor~\cite{opqccv}, meteoritic data~\cite{Dent-et-al,opqccv}, the microwave background~\cite{landau,cmb1}, stellar evolution~\cite{Eks10} and atomic clocks~\cite{Luo11}. Concerning BBN, the effect of coupled variations has mostly focused on the binding energy, $B_D$, of deuterium~\cite{flam1,fl1,Coc07} (see also Ref.~\cite{otherbd} for related investigations). The importance of $B_D$ is easily understood by the fact that the equilibrium abundance of deuterium and the reaction rate $p(n,\gamma)$D both depend exponentially on $B_D$ and on the fact that deuterium is in a shallow bound state. Indeed, in Ref.~\cite{Coc07}, we found that even a relatively small variation in the gauge or Yukawa couplings of order of a few $\times 10^{-5}$ had a significant effect on the light element abundances. In particular, using~\cite{Coc07} \begin{equation}\label{DeltaBd3} \frac{\Delta B_D}{B_D} = -13 (1+S) \, \frac{\Delta h}{h} + 18R \, \frac{\Delta \alpha}{\alpha} \, , \end{equation} a variation in the Yukawa couplings of $2 \times 10^{-5}$ induces a relative variation in $B_D$ of about 4\%. By decreasing $B_D$, nucleosynthesis begins later at a lower temperature ultimately suppressing the \li7 abundance. It is well known in principle, that the mass gaps at $A = $5 and $A=8$, prevent the nucleosynthetic chain from extending beyond \he4. Although some \li6 and \li7 is produced, their abundances remain far below that of the lighter elements, while B, Be, and CNO isotopes are produced in even smaller amounts. The presence of these gaps is caused by the instability of \he5, \li5 and \be8 with respect to particle emission: their lifetimes are as low as a few 10$^{-22}$~s for \he5 and \li5 and $\approx10^{-16}$~s for \be8. More precisely, \he5, \li5 and \be8 are respectively unbound by 0.798, 1.69 and 0.092~MeV with respect to neutron, proton and $\alpha$ particle emission\footnote{Although \li8 and ${}^8$B also contribute to the mass gap due to their short lifetimes (on BBN timescales), they are more deeply unbound, and a far greater change in the fundamental couplings would be needed to affected their stability. We will not consider them further here.}. Variations of constants will affect the energy levels of the unbound \he5, \li5 and \be8 nuclei \cite{flam1,Ber10} and hence, the resonance energies whose contributions dominate the reaction rates. In addition, since \be8 is only slightly unbound, one can expect that for even a small change in the nuclear potential, it could become bound and may thus severely impact the results of standard BBN (SBBN), in a similar way that a bound dineutron impacts BBN abundances~\cite{bbneffect}. It has been suspected that stable \be8 would trigger the production of heavy elements in BBN, in particular that there would be significant leakage of the nucleosynthetic chain into carbon. Indeed, as we have seen previously~\cite{Eks10}, changes in the nuclear potential strongly affects the triple alpha process and as a result, strongly affects the nuclear abundances in stars. This article investigates in detail the effect of the variation of fundamental constants on the properties of the compound nuclei \he5, \li5, \be8 and \carb\ involved in the \tdn, \hdp\ and \aaag\ reactions, and their consequences on reaction rates and BBN abundances. In addition, we consider the particular case of stable \be8. In Section~\ref{sec1b}, we briefly present the microscopic cluster model used to determine resonance properties. Section~\ref{sec2} focuses on \be8 and on the CNO production for the cases of both unbound and bound \be8. We compute the C/H ratio as a function of the parameter $\delta_{NN}$. Section~\ref{sec3} focuses on the \he4 production by the \hdp\ and \tdn\ reactions. Each of these reactions contains a broad s-wave resonance at low energies, and their reaction rates may depend on the resonance energies. Section~\ref{sec4} summarizes the BBN constraints on the variation of the nuclear interaction, hence extending our previous analysis~\cite{Coc07}. Section~\ref{sec5} provides a summary and our conclusions. | \label{sec5} This article investigated the influence of the variation of the fundamental constants on the predictions of BBN and extended our previous analysis~\cite{Coc07}. Through our detailed modeling of the cross-sections we have shown that although the variation of the nucleon-nucleon potential can greatly affect the triple-$\alpha$ process in stars, its effect on BBN and the production of heavier elements such as CNO is minimal at best. At the temperatures, densities and timescales associated with BBN, the changes in the \aaag\ and \beac\ reaction rates are not sufficient to produce more than C/H $\sim 10^{-21}$, and is therefore typically 6 orders of magnitude smaller than standard model abundances. This conclusion holds even when including the possibility that $^8$Be can be bound. Using the variation of the fundamental constants provides a physically motivated and well-defined model to allow for stable \be8. We have also extended previous analysis by including effects involving \he5 and \li5. This allowed us to revisit the constraints obtained in Ref.~\cite{Coc07} and in particular to show that the effect on the cross-sections remain small compared to the induced variation of $B_D$. This analysis demonstrates the robustness of our previous analysis and places the understanding of the effect of a variation of fundamental constants on BBN on a safer ground. Our analysis can be compared with Ref.~\cite{Ber10} who reached the conclusion that such variations may increase the variation of \li7 and exacerbate the lithium problem. While formally correct, our results show that this required large variation of $\delta_{NN}$ is incompatible with the BBN constraints. Note also that Ref.~\cite{Che11} assumed an independent variation of the energies of the resonances while our work considers the variation of the energies of these resonance that arises from the same physical origin, so that their amplitudes are correlated. Finally, we have extended our analysis to include the possibility of an extra relativistic degree of freedom. Because of the different dependencies, Y$(\delta_{NN}, N_\nu)$ and D/H$(\delta_{NN}, N_\nu)$, the limits on $\delta_{NN}$ or $\delta h/h$ are not relaxed for $N_\nu = 4$. The only possibility to reconcile \li7 in this context is a variation of $\delta_{NN} \sim -0.01$, along with the post BBN destruction of D/H. | 12 | 6 | 1206.1139 |
1206 | 1206.1686_arXiv.txt | Numerical simulations of minor mergers, typically having mass ratios greater than 3:1, predict little enhancement in the global star formation activity. However, these models also predict that the satellite galaxy is more susceptible to the effects of the interaction than the primary. We use optical integral field spectroscopy and deep optical imaging to study the NGC~7771+NGC~7770 interacting system ($\sim$ 10:1 stellar mass ratio) to test these predictions. We find that the satellite galaxy NGC~7770 is currently experiencing a galaxy-wide starburst with most of the optical light being from young and post-starburst stellar populations ($<1\,$Gyr). This galaxy lies off of the local star-forming sequence for composite galaxies with an enhanced integrated specific star formation rate. We also detect in the outskirts of NGC~7770 H$\alpha$ emitting gas filaments. This gas appears to have been stripped from one of the two galaxies and is being excited by shocks. All these results are consistent with a minor-merger induced episode(s) of star formation in NGC~7770 after the first close passage. Such effects are not observed on the primary galaxy NGC~7771. | Minor mergers are defined to have mass ratios greater than approximately 3:1. They have been increasingly recognized as important players in galaxy evolution and, in particular, in the formation and assembly of bulges especially in lower mass systems \citep[see][and references therein]{Hopkins2010}. Numerical simulations of minor mergers indicate that they can trigger nuclear activity and transform the morphologies of galaxies \citep{Mihos1994, Hernquist1995, Naab2003, Robertson2006, Qu2011}. The simulations of \citet{Cox2008} showed in detail that during the interaction process of minor mergers the global star formation rate (SFR) only increases moderately and the increase is a function of the mass ratio of the galaxies. Furthermore, the fractional SFR enhancement is much higher in the satellite galaxy as it is more susceptible to the tidal forces induced by the interaction. In this work we use optical imaging and integral field spectroscopy (IFS) to study the impact of the minor merger on the star formation properties of the individual galaxies of the NGC~7770+NGC~7771 system. We use these observations to test the predictions of minor merger simulations. The primary galaxy NGC~7771 is an SB(s)a galaxy, the satellite galaxy NGC~7770 is classified as an S0/a (see Fig.~\ref{fig:NOT_RGB}), and they have an approximate stellar mass ratio of 10:1 \citep{PereiraSantaella2011}. They are in a group with NGC~7769 to the northwest of the system and a small galaxy to the west of NGC~7771, with all of them being embedded in a common neutral hydrogen envelope \citep{Nordgren1997}. Throughout this work we assume a common distance of 60\,Mpc. \begin{figure} \hspace{0.7cm} \resizebox{0.8\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure1.ps}}} \caption[]{ALFOSC false-color RGB image of the NGC~7771+NGC~7700 system constructed using the SDSS broad-band $g$ and $r$ images in blue and green, respectively, and the narrow-band H$\alpha$+[N\,{\sc ii}] image in red. Orientation is north up, east to the left. The approximate FoV is $210\arcsec \times 200\arcsec$. We also mark some morphological features discussed in Section~\ref{sec:gas}.} \label{fig:NOT_RGB} \end{figure} \section[]{Observations}\label{sec:observations} We used the Potsdam Multi Aperture Spectrograph \citep[PMAS, ][]{Roth2005} in the PMAS fiber Package mode \citep[PPAK, ][]{Kelz2006} on the 3.5m telescope in Calar Alto to observe the NGC~7771+NGC~7770 system. This was part of the PPAK IFS Nearby Galaxies Survey (PINGS). Very briefly, we used the V300 grating to cover the $3700-7100$\AA \, spectral range with a spectral resolution of $10$\AA. We took three pointings with a dithered pattern, which allowed us to resample the PPAK 2.7\arcsec-diameter fiber to a final mosaic with a 2\arcsec \, spaxel and a field of view (FoV) of $148\arcsec \times 130\arcsec$. We reduced the data following \cite{Sanchez2006}. Before constructing the maps of the emission lines we fitted the stellar continuum with the {\sc fit3d} routine \citep{Sanchez2007}. We used the \cite{Bruzual2003} models to generate single stellar populations (SSP) with a Salpeter IMF covering a range of ages and metallicities. We refer the reader to \citet{Sanchez2011} for more details. We finally subtracted the fitted stellar continuum and used the observed H$\alpha$/H$\beta$ line ratio to correct for extinction the line maps on a spaxel-by-spaxel basis \citep[see][for a full description]{RosalesOrtega2010}. We also obtained imaging with the Andalucia Faint Object Spectrograph and Camera (ALFOSC) on the 2.5m Nordic Optical Telescope (NOT) in El Roque de los Muchachos. We used the SDSS broad-band $g$ and $r$ filters and an H$\alpha$+[N\,{\sc ii}] narrow-band filter ($\lambda_{\rm c}= 6653$\AA \, and $\Delta \lambda=55$\AA \, at full width half maximum FWHM). The instrument has a 0.19\arcsec/pixel plate scale and the observations were obtained under photometric and extremely good seeing ($\sim 0.6\arcsec$, FWHM as measured from stars in the FoV) conditions. We followed standard techniques to reduce and calibrate the data, as well as to subtract the continuum from the H$\alpha$+[N\,{\sc ii}] image and compute the H$\alpha$ emission. \section[]{Results}\label{sec:results} \subsection{Optical continuum and ionized gas properties}\label{sec:gas} The deep ALFOSC optical continuum images allow us to get a detailed view of the morphology of the system, especially of the lesser known galaxy NGC~7770. Although NGC~7770 has been classified as an S0/a, it shows a faint spiral structure with bright knots located in a ring of star formation (Fig.~\ref{fig:NOT_RGB}). There are also other interesting features seen in the optical continuum emission such as a truncated stellar disk in NGC~7771 and faint short tidal tails in both galaxies. Such continuum tail features are predicted by simulations of minor mergers right after the first passage \citep[][]{Cox2008}. Both galaxies in the system show extended H$\alpha$ emission over several kpc (Fig.~\ref{fig:PPAK_maps}). In NGC~7771 the H$\alpha$ emission arises from the nuclear region of the galaxy, which is in fact a ring of star formation \citep[see][and references therein]{AAH10}, from bright H\,{\sc ii} regions at the end of the bar and in the spiral arms, and from the disk of the galaxy. NGC~7770 also shows a circumnuclear ring of star formation with an approximate diameter of 20\arcsec \, ($\sim 5.7\,$kpc). The rings of star formation in both galaxies show optical line ratios typical of H\,{\sc ii} regions, as can be seen from the PPAK spectral maps (Fig.~\ref{fig:PPAK_maps}) and the spatially-resolved diagnostic diagrams (Fig.~\ref{fig:PPAK_BPTs}). For the spaxels in both galaxies showing H\,{\sc ii} region-like line ratios, the differences in the observed [O\,{\sc iii}]/H$\beta$ line ratios are readily explained by differences in the gas-phase oxygen abundances (see Section~\ref{sec:abundances}). \begin{figure*} \hspace{0.5cm} \resizebox{0.32\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure2a.ps}}} \resizebox{0.32\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure2b.ps}}} \hspace{0.5cm} \resizebox{0.32\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure2c.ps}}} \resizebox{0.32\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure2d.ps}}} \caption[]{PPAK spectral maps (color images) of the observed H$\alpha$ emission (units of $10^{-16}\,{\rm erg\,cm}^{-2}\,{\rm s}^{-1}\,{\rm arcsec}^{-2}$) and the extinction-corrected [O\,{\sc iii}]$\lambda$5007/H$\beta$, [N\,{\sc ii}]$\lambda$6583/H$\alpha$, and [S\,{\sc ii}]$\lambda\lambda$6717,6731/H$\alpha$ line ratios. The contours (in a square root scale) are the ALFOSC continuum-subtracted H$\alpha$ emission.} \label{fig:PPAK_maps} \end{figure*} \begin{figure*} \resizebox{0.65\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure3a.ps}}} \vspace{-3.5cm} \resizebox{0.65\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure3b.ps}}} \vspace{-3.7cm} \caption[]{Spatially-resolved diagnostic diagrams. Each measurement corresponds to one spaxel. The star-like symbols are the nuclear line ratios, and the small and large dots are spaxels at projected galactocentric distances of $R<12\arcsec$ and $R>12\arcsec$, respectively. The solid curves are ``maximum starburst lines'' \citep{Kewley2001} and the dashed curves the empirical separation between Composites and H\,{\sc ii} regions and Seyferts and LINERs \citep{Kauffmann2003BPT, Kewley2006}. The spaxels of NGC~7770 (upper panel) in the LINER/Seyfert and AGN regions can also be produced by shocks \citep[see text and also][]{Allen2008}. }\label{fig:PPAK_BPTs} \end{figure*} The most striking feature are the filaments of diffuse H$\alpha$ emission in the interface region between the two galaxies to the northwest of NGC~7770, and also to the south of the galaxy (see Fig.~\ref{fig:NOT_RGB}). These H$\alpha$ filaments are not detected in the optical continuum emission (Fig.~\ref{fig:NOT_RGB}) indicating that they are mostly made of gas. The northwest filaments extend out to projected galactocentric distances of approximately 6\,kpc. Minor-merger simulations \citep[][]{Cox2008} predict that gas is likely to have been cleared out from the primary galaxy extended disk after the first close passage of the satellite galaxy. The H$\alpha$ filaments of NGC~7770 and a few spaxels tracing the diffuse H$\alpha$ emission around the ring of star formation of NGC~7771 \citep[see also][]{AAH10} display elevated optical line ratios. This is better seen in the spatially-resolved diagnostic diagrams (Fig.~\ref{fig:PPAK_BPTs}) where a large fraction of the outer (i.e., projected galactocentric distances $>12\arcsec \sim 3.5\,$kpc) H$\alpha$ emitting regions of NGC~7770 are in the composite/AGN region or the LINER/Seyfert region. The observed line ratios for the filament spaxels, even those falling in the LINER/Seyfert region, can be reproduced with the shock+precursor models of \citet{Allen2008} with shock velocities of $300-400\,{\rm km\,s}^{-1}$ (their figures 31, 32, and 33). This could be understood if gas, from either NGC~7771 or NGC~7770, is being stripped and excited by shocks induced by tidal forces. In NGC~7771 there is no variation of the optical line ratios with the galactocentric distance. \subsection{Luminosity-weighted stellar ages and metallicities}\label{sec:ages} The luminosity-weighted ages and metallicities of stellar populations can give us hints about the {\it recent} star formation history of galaxies. To ensure an appropriate modelling of the stellar populations we used a Voronoi binning method \citep{Cappellari2003} and the {\sc pingsoft} \citep{RosalesOrtega2011} software to define regions (also known as voxels) in the PPAK data cube with extracted spectra of sufficient signal-to-noise (S/N) ratios. We targeted S/N ratios of 25 over a spectral region centered at 5200\,\AA \, with a 100\,\AA-width. We then used the {\sc fit3d} routine to fit SSPs to the entire spectrum of each voxel with a $\chi^2$ minimization technique. We generated a grid of templates using the \citet{Bruzual2003} models with a range of ages (5\,Myr, 25\,Myr, 100\,Myr, 640\,Myr, 1.4\,Gyr, 2.5\,Gyr, 5\,Gyr, 13 Gyr, and 17 Gyr) and 3 metallicities ($Z=0.008$, 0.02, and 0.05), and allowed for the templates to be attenuated. In general it was not possible to distinguish among the few best-fitting models. We therefore estimated the age and metallicity of each extracted spectrum by averaging the ages and metallicities of these models, weighted by their corresponding $\chi^2$ values. For the following discussion we grouped the stellar populations in young and post-starburst with ages of $<$1\,Gyr, intermediate with ages of $1-5\,$Gyr, and old with ages $>$5\,Gyr. The optical light of the primary galaxy is mostly dominated by intermediate and old stellar populations (see Fig.~\ref{fig:PPAKNOT_stellar}) with super-solar metallicities ($Z>0.025$). There is a small contribution ($\sim 20\%$) from stellar populations with ages of $<$1\,Gyr, which are mostly located in the circumnuclear ring of star formation. This agrees with the presence of post-starburst stellar populations in the nuclear region of NGC~7771 found by previous works \citep[see][]{Davies1997, AAH10}. In the satellite galaxy NGC~7770 the young stellar populations ($<1$Gyr) with solar and super-solar metallicities ($Z=0.015-0.04$) contribute as much as $\sim 80$\% of the optical light. The presence of a post-starburst stellar population has often been interpreted as the result of a past interaction \cite[e.g.,][]{Liu1995}. The fact that the optical emission of NGC~7770 is dominated by this population suggests that the galaxies have already experienced first passage. This likely triggered a strong burst of star formation in the satellite galaxy less than 1\,Gyr ago, as predicted by the simulations of \citet{Cox2008}. The origin of the post-starburst population in NGC~7771 is less clear since these minor-merger models do not predict a strong star formation enhancement in the primary. Based on the neutral hydrogen morphology, \citet{Nordgren1997} suggested that NGC~7771 and NGC~7769, which have a 2:1 mass ratio, appear to be having a prograde--retrograde interaction, with NGC7769 being the retrograde one. Such an interaction, while not as efficient at triggering strong episodes of star formation, may explain the circumnuclear burst of star formation in NGC~7771. \subsection{Current Star Formation Activity}\label{sec:sfr} Next we explore the impact of the interaction on the current SFR of the galaxies. For a 10:1 stellar mass ratio \citet[][]{Cox2008} predict only a mild total enhancement ($< 2 \times$) in the global SFR of the system when compared to the sum of SFRs of the individual galaxies if they were isolated. In fact, most of the SFR enhancement takes place in the satellite galaxy because it is that experiencing significant tidal forces. We can test this prediction for the individual galaxies of this system since in the local Universe star-forming galaxies define a sequence of increasing SFR for increasing stellar masses \citep[e.g.,][and references therein]{Salim2007}. \begin{figure} \hspace{0.95cm} \resizebox{0.7\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure4a.ps}}} \vspace{-1.2cm} \hspace{0.95cm} \resizebox{0.7\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure4b.ps}}} \vspace{-1.3cm} \caption[]{Distributions of the ages and metallicities of the spatially-resolved (i.e., corresponding to the extracted spectra for the defined voxels, see text) modelled SSPs. The colors represent the fractional contribution of the SSP bins (luminosity-weighted) to the observed optical emission ($3700-7100\,$\AA) of each galaxy. } \label{fig:PPAKNOT_stellar} \end{figure} The integrated extinctions $A_V$ of the galaxies ($1.1\,$mag and $2.5\,$mag for NGC~7770 and NGC~7771, respectively, from the H$\alpha$/H$\beta$ ratio) indicate that we can use the SFR recipe of \citet{Kennicutt2009}. It combines the observed (not corrected for $A_V$) H$\alpha$ luminosity and the IR luminosity to account for the unobscured and the obscured SFRs, respectively. The observed H$\alpha$ luminosities are $L({\rm H}\alpha) = 2.7 \times 10^{41}\,{\rm erg \,s}^{-1}$ and $4.2 \times 10^{41}\,{\rm erg \,s}^{-1}$ for NGC~7770 and NGC~7771, respectively, and the IR luminosities are $\log (L_{\rm IR}/{\rm L}_\odot) = 10.8$ and $11.3$ \citep[][]{PereiraSantaella2011}. There is some evidence of the presence of an AGN in NGC~7770 both from the composite nuclear (i.e., star-forming/AGN) [N\,{\sc ii}]/H$\alpha$ line ratio (Fig.~\ref{fig:PPAK_BPTs}) and the hot dust emission detected in the Spitzer spectrum \citep{AAH12}. To compute the SFR we then need to subtract the AGN contribution from the IR luminosity. For a Kroupa IMF we obtained SFR(total)$= 4.5\,{\rm M}_\odot \,{\rm yr}^{-1}$ and $12.8\,{\rm M}_\odot \,{\rm yr}^{-1}$ for NGC~7770 and NGC~7771, respectively. These are similar, within the uncertainties of the SFR recipes ($\pm0.3\,$dex or better), to those obtained from the extinction-corrected L(H$\alpha$). We note that the SFR recipe used here assumes that approximately half of the dust heating arises from stars older than 30\,Myr. The IR-based calibration of \citet{Kennicutt1998}, which assumes that the dust is only heated by these young stars, results in SFRs approximately two times higher. The specific SFR (SSFR$={\rm SFR}/M_*$) using stellar masses computed with a near-IR mass-to-light ratio \citep[][]{Bell2001, PereiraSantaella2011} are $\log ({\rm SSFR/yr}^{-1})=-9.6$ and $-10.4$ for NGC~7770 and NGC~7771, respectively. If there was no AGN in NGC~7770 then the calculated SFR and SSFR would be lower limits. We find that NGC~7771 falls on the high mass end of the local star-forming sequence. Composite galaxies and AGN follow a similar star-forming sequence but extending to higher stellar masses \citep[][their figure~18]{Salim2007}. The satellite galaxy NGC~7770 shows an elevated SSFR and lies just off of the distinct region occupied by local composite galaxies. This suggests that not only does NGC~7770 show evidence of a post-starburst stellar population, but that it is also currently undergoing a galaxy-wide burst of star formation. These results agree with predictions of minor merger simulations as well as with statistical observational studies of minor mergers \citep{Woods2007,Ellison2008}. \subsection{Gas-phase oxygen abundances}\label{sec:abundances} Interacting galaxies are believed to contribute to the observed scatter in the local mass-metallicity relation \citep{Tremonti2004}. Specifically, local luminous and ultraluminous IR galaxies, which include a large fraction of gas-rich mergers, show lower nuclear abundances than local emission-line galaxies of similar luminosity and mass \citep{Rupke2008}. Simulations of interacting galaxies show that it is possible to decrease the nuclear metallicity although the gas content of the galaxies plays a role \citep{Torrey2012}. We used the calibration of \citet{Pettini2004} based on the O3N2 index to derive the oxygen abundances of NGC~7770 and NGC~7771 on a spaxel-by-spaxel basis, but only for regions with [O\,{\sc i}]/H$\alpha$$<$0.06. In doing so we ensured that most of the line emission used to compute the abundances is produced by photoionization in H\,{\sc ii} regions rather than by other mechanisms. The oxygen abundances of the regions of NGC~7770 are lower that those of NGC~7771 (Fig.~\ref{fig:metallicities}) as expected from the different stellar masses and in agreement with the SSP modelling. The distribution of the oxygen abundance in NGC~7771 is rather complex with the H\,{\sc ii} regions at the east end of the bar having apparently lower abundances than those to the west. The abundance in the central regions of NGC~7770 appears to be about 0.1dex lower than in the bright H\,{\sc ii} region to the southwest of the nucleus and in the outskirts of the galaxy. This may occur if gas from the outskirts of the satellite galaxy has been channeled to the nuclear region as the result of the interaction. However the suspected composite nature of the nucleus of NGC~7770 may complicate this interpretation. | \label{sec:conclusions} We presented PPAK spatially-resolved optical emission line ratio and abundance maps and diagnostic diagrams as well as ALFOSC deep optical imaging of the NGC7771+NGC7770 system to study the impact of the interaction on the star formation properties of the individual galaxies. The galaxies have a $\sim$ 10:1 stellar mass ratio and thus are classified as a minor merger. Numerical simulations of minor mergers predict little global enhancement of the star formation properties and mostly affecting the satellite galaxy. We find that the satellite galaxy in the system NGC~7770 is indeed experiencing a galaxy-wide starburst with most of of its optical emission being produced by young and post-starburst stellar populations ($<1\,$Gyr). This galaxy shows an enhanced SSFR with respect to the distinct region occupied by composite objects in the local star-forming sequence. We also detected H$\alpha$ emitting gas filaments in the outskirts of NGC~7770 at projected galactocentric distances greater than $\sim 4\,$kpc and out to 6\,kpc. These filaments appear to be made of gas that has been stripped from one of the two galaxies and is being excited by shocks. Overall, our findings support the picture that the minor merger induced one or several episodes of star formation in the satellite galaxy NGC~7770 after the first close passage. Such interaction-induced effects are not observed in the primary galaxy of the system. We thank a referee for comments that helped improve the paper and M. Cappellari for interesting discussions. A.A.-H. acknowledges finantial support from the Spanish Plan Nacional grant AYA2010-21161-C02-01 and the Universidad de Cantabria AGL program, A.I.D. from grant AYA2010-21887-C04-03, and F.F.R.-O. from the Mexican National Council for Science and Technology (CONACYT) through the programme Estancias Posdoctorales y Sab\'aticas al Extranjero para la Consolidaci\'on de Grupos de Investigaci\'on, 2010-2011. M.P.-S. is funded by an ASI fellowship under contract I/005/11/0. \begin{figure} \hspace{0.3cm} \resizebox{0.9\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure5.ps}}} \vspace{-1.5cm} \caption[]{Map of oxygen abundances derived with the O3N2 index and the calibration of \citet{Pettini2004} only for spaxels with no evidence of shock excitation ([O\,{\sc i}]$\lambda$6300/H$\alpha <0.06$). Contours are as in Fig.~\ref{fig:PPAK_maps}.} \label{fig:metallicities} \end{figure} | 12 | 6 | 1206.1686 |
1206 | 1206.1962_arXiv.txt | Here we show an example of a young asteroid cluster located in a dynamically stable region, which was produced by partial disruption of a primitive body about 30 km in size. We estimate its age to be only $1.9\pm0.3$ Myr, thus its post-impact evolution should have been very limited. The large difference in size between the largest object and the other cluster members means that this was a cratering event. The parent body had a large orbital inclination, and was subject to collisions with typical impact speeds higher by a factor of $2$ than in the most common situations encountered in the main belt. For the first time we have at disposal the observable outcome of a very recent event to study high-speed collisions involving primitive asteroids, providing very useful constraints to numerical simulations of these events and to laboratory experiments. | \label{} The asteroid population, being steadily subject to a process of collisional evolution \citep{davis1989,bottke2005,morby2009,asphaug2009}, provides excellent possibilities to study physics of collisional events. Asteroid families, which are believed to originate from catastrophic disruption of single parent bodies \citep{Zappala2002}, are, almost one century since the pioneering work by \citet{hirayama1918}, still an attractive and challenging subject. They provide a key to our understanding of the collisional history of the main asteroid belt \citep{bottke2005,cellino2009}, outcomes of disruption events over a size range inaccessible to laboratory experiments \citep{michel2003,durda2007,asphaug2010}, clues on the mineralogical structure of their parent bodies \citep{cellino2002}, the role of space weathering effects \citep{Nesvorny2005,vernazza2009} and to many other subjects. So far, ejecta from a few tens of large-scale collisions has been discovered across the main asteroid belt \citep[e.g.][]{Zappala1995,MD2005,Nesvorny2005}. In terms of their estimated ages, most families identified so far are fairly old and have had enough time to evolve significantly since the epoch of their formation as a consequence of (i) chaotic diffusion \citep{nes2002flora,Novakovic2010b}, (ii) semi-major axis drift due to Yarkovsky effect \citep{farinella1999,bottke2001}, (iii) secondary collisions \citep{Marzari1999,bottke2005}, (iv) non-destructive collisions \citep{aldo2007} and/or (v) diffusion due to close encounters with massive asteroids \citep{carruba2003,Novakovic2010c,Delisle2012}. In this respect, little altered recently born families may provide more direct information about the physics of break-up events. Evidence of recent collisions in the asteroid belt have been reported in the last decade and our knowledge about young asteroid families has been increased significantly \citep{Nesvorny2002,datura2006,nes2006}. Most of these groups are formed by asteroids belonging to the $S$ taxonomic class. There are, however, several important differences among the $S$ and $C$-type asteroids. The objects belonging to former class are thought to have experienced some thermal evolution since the time of their formation, and it is, for example, known that space weathering processes are different for these two classes of objects \citep[e.g.][]{gaffey2010}. Also, numerical simulations show that the outcomes of collisional events are dependent on internal structure of the parent body \citep{jutzi2009}. Because of these reasons it is necessary to identify also young $C$-class families in dynamically stable regions, because a few such groups are already known, but none of these is well suited to extract reliable enough information. Two $C$-type families, namely Veritas and Theobalda, about 8.3 and 6.7 Myr old respectively, are both located in dynamically unstable region \citep{nes2003,Novakovic2010a}. Thus, despite their young ages, these families evolved significantly since post-impact situation. Most of the asteroids belonging to Beagle family \citep{Nesvorny2008}, which is probably less than 10 Myr old, are located in dynamically relatively stable region. However, this group is embedded in the large Themis family making distinction between the real members of the group and background objects very difficult. Finally, the youngest known group that might be formed by $C$-type asteroids is Emilkowalski cluster, which is only $220\pm30$ kyr old \citep{nes2006}. However, it seems to be rather an $X$- than $C$-type group because albedos of its members are much higher than expected for $C$-type objects. For example, geometric albedo of asteroid (14627) Emilkowalski is $0.2013\pm0.0170$ \citep{wise}. Thus, it is of extreme importance to identify young families, that belong to the most primitive $C$ class, that do not suffer from above mentioned problems. We have found the first example of this kind to be the Lorre cluster, recently discovered by \citet{Novakovic2011}. According to existing color data its largest member, (5438)~Lorre, is a primitive carbonaceous $C$-class asteroid, which may contain organic materials. Moreover, the members of this cluster are located in dynamical stable region and very tightly packed in the space of proper orbital elements \citep{Kne2003}, suggesting a likely young age. Therefore, its post-impact evolution should have been very limited. This makes it a very promising candidate for different possible studies. Two crucial prerequisites for these studies are an accurate identification of its members, and a reliable estimation of its age. These are the questions we address here. | \label{s:conclusions} Here we show the first example of a young asteroid cluster located in a dynamically stable region, which was produced by partial disruption of a primitive body about 30 km in size. We estimate its age to be only $1.9\pm0.3$ Myr, thus its post-impact evolution is very limited. The large difference in size between the largest object and the other cluster members means that this was a cratering event. The parent body had a large orbital inclination, and was subject to collisions with typical impact speeds higher by a factor of $2$ than in the most common situations encountered in the main belt. For the first time we have at disposal the observable outcomes of a very recent event to study high-speed collisions involving primitive asteroids, providing very useful constraints to numerical simulations of these events \citep{michel2003,jutzi2009,leinhardt2012} and to laboratory experiments \citep{housen2011}. This is the best preserved young asteroid family produced by partial disruption of a primitive asteroid, of a kind which is supposed to have survived nearly unaltered since the epoch of formation of the Solar System. Being young and well distinct from the background population, this cluster provides very useful information that can help to answer several long-debated questions in planetary science. Examples include a better understanding of impact physics, material strength and the role of space weathering. These process, highly dependent on the composition of the objects, are so far poorly constrained for primitive asteroids. Among the members of the Lorre cluster there are several asteroid pairs, couples of objects with nearly identical orbital parameters. These pairs may well consist of couples of fragments which were ejected with nearly identical ejection velocities. Another possibility is that they might actually be the components of former binary systems originally produced by the collision, and later decoupled by some mechanisms \citep{pravec2010}. Production of binary systems in collisional events has been suggested by numerical simulations \citep{michel2001,durda2004}, but their expected abundance in asteroid families has not been firmly established yet. The young age of the Lorre cluster as well as its sharp separation from background objects may potentially help to better understand both populations, binaries and pairs. An interesting possibility for future work comes from a recent result of \citet{benavidez2012} how found that low-energy impacts into rubble-pile and monolithic targets produce different features in the resulting SFD, and, thus, this is a potentially diagnostic tool to study the initial conditions just after the impact and the internal structure of the parent bodies of asteroid families. According to \citet{benavidez2012}, cratering events, produced by small impactors, can potentially provide even more information about the internal structure of the parent body than catastrophic or super-catastrophic events produced by large impactors. Thus, the Lorre cluster seems to be a very promising candidate. Next, the Lorre cluster may be very useful to improve our knowledge about space weathering processes acting on primitive bodies, a debated subject since results based on the Sloan Digital Sky Survey broadband photometry \citep{Nesvorny2005} are not consistent with the results of some laboratory experiments \citep{brunetto2009}. Finally, the cluster may be a very interesting place to search for new main-belt comets (MBCs). \footnote{Main belt comets are objects dynamically indistinguishable from main belt asteroids, but which exhibit comet-like activity due to the sublimation of volatile ice \citep{Hsieh2006}.} A recent findings by \citet{nov2012} supports an idea that this kind of objects may be preferentially found among the members of young asteroid families \citep{Nesvorny2008,Hsieh2009}. In this respect, members of Lorre cluster are particularly interesting candidates because their heliocentric distances are smaller than those of currently known MBCs. Thus, they may provide a clue about the inner edge of populations of MBCs. | 12 | 6 | 1206.1962 |
1206 | 1206.6897_arXiv.txt | { We present supernova rate measurements at redshift 0.1--1.0 from the Stockholm VIMOS Supernova Survey (SVISS). The sample contains 16 supernovae in total. The discovered supernovae have been classified as core collapse or type Ia supernovae (9 and 7, respectively) based on their light curves, colour evolution and host galaxy photometric redshift. The rates we find for the core collapse supernovae are $3.29_{-1.78\,-1.45}^{+3.08\,+1.98}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ (with statistical and systematic errors respectively) at average redshift 0.39 and $6.40_{-3.12\,-2.11}^{+5.30\,+3.65}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ at average redshift 0.73. For the type Ia supernovae we find a rate of 1.29$_{-0.57\,-0.28}^{+0.88\,+0.27}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ at $\langle z\rangle =0.62$. All of these rate estimates have been corrected for host galaxy extinction, using a method that includes supernovae missed in infrared bright galaxies at high redshift. We use Monte Carlo simulations to make a thorough study of the systematic effects from assumptions made when calculating the rates and find that the most important errors come from misclassification, the assumed mix of faint and bright supernova types and uncertainties in the extinction correction. We compare our rates to other observations and to the predicted rates for core collapse and type Ia supernovae based on the star formation history and different models of the delay time distribution. Overall, our measurements, when taking the effects of extinction into account, agree quite well with the predictions and earlier results. Our results highlight the importance of understanding the role of systematic effects, and dust extinction in particular, when trying to estimate the rates of supernovae at moderate to high redshift.} | The cosmic rate of supernovae is an important observable that can be used to constrain the properties of galaxies at high redshifts and to study the supernovae (SNe) themselves. Depending on the origin of the SN explosion, thermonuclear or core collapse, different aspects can be studied. The first measurement of the cosmic SN rate (SNR) was done by \citet{1938ApJ....88..529Z} who found that ``the average frequency of occurrence of supernovae is about one supernova per extra-galactic nebula per six hundred years'' for the local volume. It is not until the latest decades that the higher redshift regimes have been possible to reach. More recently, large programmes targeting type Ia supernovae at intermediate and high ($\gtrsim 0.1$) redshifts have been conducted to measure the expansion of the universe and do precision cosmology \citep[e.g. ][]{1997ApJ...483..565P,1998ApJ...507...46S, 2006A&A...447...31A,2007ApJ...659...98R,2007ApJ...666..674M}. Some of these surveys also report supernova rates for Ia SNe \citep[e.g.][]{2006AJ....132.1126N,2008ApJ...681..462D} out to $z\sim 1.5$ and core collapse supernovae (CC SNe) \citep{2004ApJ...613..189D,2009A&A...499..653B} out to $z\sim 0.7$. Large surveys targeting SNe of any type have also been successful in finding and characterising supernovae as well as determining both local and intermediate redshift cosmic supernova rates \citep{2008A&A...479...49B,2010ApJ...713.1026D, 2011MNRAS.tmp..413L}. \citet{2010ApJ...718..876S}, \citet{2010ApJ...715.1021D}, and \citet{2012ApJ...745...32B} survey galaxy clusters, where the SN Ia rates are likely to be enhanced, to find supernovae and have reported cluster SN rates out to redshift 0.9. \citet{2012ApJ...745...31B} also report SN Ia rates from detections in the foreground and background of the targeted galaxy clusters out to $z\sim 1.5$. The standard observational strategy for finding SNe at high redshift is to perform survey observations on a given field and then monitor the same field over a long period of time. Supernovae are discovered by searching the images for variable sources using image subtraction tools \citep[such as ][]{2000A&AS..144..363A} to minimise subtraction residuals. The cadence of observations during the survey period is normally chosen to sample signature features of the SN light curves and colour evolution at the target redshifts. In this way photometric typing of the SNe is possible and the light curves can be used to study the SN characteristics. For the Ia surveys with cosmology as the main goal, follow-up spectroscopy of SN Ia candidates is needed to get a secure determination of the redshift, to improve the accuracy in the distance measurement, and confirmation of the type. When calculating supernova rates from this kind of survey, care has to be taken to avoid selection effects from the spectroscopic observations. Supernova typing normally includes studying the spectra of the SNe close to their peak luminosity and identifying spectral lines, notably H, He and Si\rm{II} lines, something which is observationally very expensive and in practice unfeasible at high redshift for fainter SN types. Another method to type SNe is to compare the observed light curves and colour evolution to pre-existing templates of different SN types, i.e. photometric SN typing. These methods have been demonstrated to work \citep[e.g. ][]{2007ApJ...659..530K,2007AJ....134.1285P,2009ApJ...707.1064R, 2010PASP..122.1415K} using somewhat different techniques. Better typing accuracy is achieved with prior information on the redshift through photometric or spectroscopic redshift of the host galaxy \citep[][]{2010PASP..122.1415K,2011A&A...532A..29M}. Thermonuclear supernovae (or SN Ia's) are thought to be white dwarfs that explode when they accrete matter and approach the Chandrasekhar limit \citep[for a review, see][]{2000A&ARv..10..179L}. When taking the luminosity-stretch relation \citep{1993ApJ...413L.105P} into account the peak luminosity of these SNe exhibit a very narrow spread and can thus be used to accurately measure cosmological distances. The exact details of the explosions and of the progenitor system are not fully understood. For example, the time between formation of the progenitor system and the supernova explosion -- the so called delay time -- is unknown. This delay time depends on the nature of the companion star to the white dwarf \citep{2005A&A...441.1055G}. By studying the rates of Ia supernovae and comparing to either the cosmic star formation history \citep[e.g.][]{1999A&A...350..349D, 2006AJ....132.1126N,2010ApJ...713...32S}, or the star formation rates and stellar masses of the host galaxies \citep{2006ApJ...648..868S,2008PASJ...60.1327T,2011MNRAS.tmp..307M} it is possible to set constraints on the delay time and thereby also on the progenitor system. Core collapse supernova explosions (CC SNe) are the end-points of the lives of massive stars, with masses between $\sim 8 \,\mbox{M}_{\sun}$ and $\sim 50 \, \mbox{M}_{\sun}$ \citep{1984ApJ...277..791N, 1997ApJ...483..228T,2009ARA&A..47...63S}. Since massive stars are short-lived compared to the cosmic time-scales the CC SNe trace active star formation. By averaging the CC SN rate over cosmic volume the rest frame rate of star formation in that volume can be studied. In this way an independent measure of the star formation history at high redshift can be obtained \citep{2004ApJ...613..189D,2005A&A...430...83C, 2008A&A...479...49B, 2009A&A...499..653B}. More conventional methods of finding the cosmic star formation rates include measuring the rest-frame UV light from galaxies at a given redshift \citep[e.g. ][]{2004ApJ...600L.103G,2009ApJ...705..936B}, measuring the far-infrared (FIR) light to take star formation hidden by dust into account \citep[e.g.][]{2005ApJ...632..169L} and deriving the rates from H$\alpha$ measurements \citep{2009ApJ...696..785S,2010A&A...509L...5H}. The UV and H$\alpha$-based methods have a drawback, that is also present for the SNR method, in that a correction for dust extinction needs to be applied. Methods based on using the FIR light to estimate the added star formation from re-radiated UV light make it possible to correct the star formation history for dust extinction effects. \citet{2006ApJ...651..142H} presented a compilation of the star formation history from multiple sources for $z\sim 0-6$, taking the obscured star formation into account, and fitted an analytical function to the data. The light from a supernova has to pass through its host galaxy before starting on the long trip to reach our telescopes. When travelling through the gas and dust inside its host the supernova light will be subject to varying degrees of extinction, depending on the dust content of the galaxy and the position of the SN with respect to the observer (e.g. a SN situated in an edge-on galaxy will suffer from higher extinction, on average, than one in a face-on galaxy). \citet{1998ApJ...502..177H} and \citet{2005MNRAS.362..671R} present models of extinction for core collapse and thermonuclear SNe in normal spiral galaxies. These models are built by using Monte Carlo simulations of supernova positions within a galaxy with given morphology and dust content. By using the extinction models it is possible to estimate the effect on the observed supernova rates \citep{2004ApJ...613..189D}. It should be noted that this method is mainly applicable to normal galaxies with low to medium amounts of gas/dust, typical of galaxies in the local volume of the universe. As the redshift increases, more and more of the star formation takes place inside dusty galaxies. \citet{2005ApJ...632..169L, 2009A&A...496...57M, 2011A&A...528A..35M} find that the star formation from these infrared bright galaxies dominate the global star formation at redshift 0.7 and higher. In these galaxies the SN explosions can be completely obscured by the large amounts of dust within the nuclear regions. For low to moderate amounts of extinction this can be estimated and taken into account in the light curve analysis. But for host galaxies with high dust content (such as luminous and ultra-luminous infrared galaxies, LIRGs and ULIRGs) most of the SNe may not be detectable, even in the near-infrared (NIR) where the extinction is strongly reduced \citep[e.g.][]{2001MNRAS.324..325M}. When calculating the rates, the number of missing SNe due to the change in average extinction in star forming galaxies with redshift needs to be compensated for \citep{2007MNRAS.377.1229M}. The derivation of the de-bias factors that are needed to correct the rates for this effect is further complicated by the recent discovery that the population of U/LIRGs at low redshift is quite different from the ones at high redshift \citep[e.g.][]{2010ApJ...713..686D, 2011A&A...533A.119E, 2011arXiv1110.4057K}. \citet{2012arXiv1206.1314M} have recently estimated the fraction of SNe missed in such galaxies both locally and as a function of redshift making use of the most recent results on the nature of U/LIRGs at different redshifts. The Stockholm VIMOS Supernova Survey (SVISS) is a multi-band ($R+I$) imaging survey aimed at detecting supernovae at redshift $\sim$0.5 and derive thermonuclear and core collapse supernova rates. The supernova survey data were obtained over a six month period with VIMOS/VLT. \citet{2008A&A...490..419M} describe the supernova search method along with extensive testing of the image subtraction, supernova detection and photometry. The discovery and typing of 16 supernovae in one of the search fields is reported in \citet{2011A&A...532A..29M}. Here we present the supernova rates estimated from the survey data along with delay time distributions for the Ia SNe and star formation rates (SFR) calculated from the CC SNe rates. The first part of the paper contains a description of the observations and supernova sample. In Sect.~\ref{sec:method} we describe the method used to determine the supernova rates and in Sect.~\ref{sec:results} the resulting supernova rates, delay time distribution and star formation rates are presented. In the final section of the paper, Sect.~\ref{sec:disc}, we discuss the results and compare them to the work of other authors. The Vega magnitude system and a standard $\Lambda$CDM cosmology with $\{H_0,\Omega_{M},\Omega_{\Lambda}\}=\{70,0.3,0.7\}$ have been used throughout the paper. | \label{sec:disc} We have presented supernova rates from the SVISS along with a description of the methods used to compute them and an extensive analysis of systematic errors for the rates. The resulting rates for the core collapse SNe are: $3.29_{-1.78\,-1.45}^{+3.08\,+1.98}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ (with statistical and systematic errors respectively) at $\langle z\rangle =0.39$ and $6.40_{-3.12\,-2.11}^{+5.30\,+3.65}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ at $\langle z\rangle=0.73$. The CC rates have been corrected for obscuration in dusty environments using the results of \citet{2012arXiv1206.1314M}, as described in Sect.~\ref{sec:obsc} and host galaxy extinction based on the models in \citet{2005MNRAS.362..671R}. Uncorrected values can be found in Table~\ref{table:rates}. The rate estimates follow the star formation history well. \citet{2011ApJ...738..154H} point out that the core collapse supernova rates found in both local and high redshift searches seem to be too low by a factor of two when compared to the star formation history compiled in \citet{2006ApJ...651..142H}. Our results do not show this difference, both of our rate estimates are consistent with the star formation history within the statistical errors. But our errors are quite large, the rates estimated by other authors, in particular at low $z$, do in fact differ significantly from the SFH. The strongest evidence for this comes from the rate estimate by \citet{2011MNRAS.tmp..317L} at low $z$ which has small errors and is clearly below the SFH. At higher redshift the problem is less severe, which could be due to the increased statistical errors. At high redshift the problem may also be somewhat alleviated by the obscuration corrections, which we have included in our plotted rates in Fig.~\ref{fig:ccrates}, the rates from other surveys plotted in this figure do not include this correction (with the exception of the data point from \citealt{2011MNRAS.tmp.1508G}). \citet{2011ApJ...738..154H} suggest that taking missed SNe due to extinction and dust obscuration in LIRGs and ULIRGs into account is not enough to explain the difference. Instead they suggest that the reason is that the assumed fractions of faint and very faint CC SNe are too low. In our tests of the systematic uncertainties we find that assuming 30\% of the CC SNe to be faint ($M_B>-15$) boosts the rates by $\sim$30\%, not enough to bridge the factor of two gap found for the local SN searches. Of course, the assumption on fractions of faint CC SNe may be very different for the different surveys, making it possible that other data points may go up more than this. While the SFHs we compare with have also been corrected for dust and hidden star formation, the supernovae are probably affected differently by the presence of obscuring dust. The number of supernovae that are missed in dusty starburst galaxies and LIRGs is currently not well constrained even in the local universe \citep[e.g][]{2003A&A...401..519M, 2004NewAR..48..595M}. The de-bias factors for extinction and dust obscuration in normal galaxies and U/LIRGs derived by \citet{2012arXiv1206.1314M} can make the difference between the predicted and observed CC SN rates disappear. In this paper we have used these factors to de-bias the core collapse supernova rates. The de-bias factors from this study are slightly larger at low redshifts, but lower at high redshift, than the factors from \citet{2007MNRAS.377.1229M}. The uncertainties -- both statistical and systematic -- of the de-bias factors derived in \citet{2012arXiv1206.1314M} have been thoroughly studied, and will hopefully decrease as more observations of SNe in dusty galaxies are obtained. High angular resolution observations at near-IR \citep[e.g.][]{2007ApJ...659L...9M, 2008ApJ...689L..97K, 2012ApJ...744L..19K} and radio \citep[e.g,][]{2010A&A...519L...5P, 2011MNRAS.415.2688R} wavelengths have recently been used to detect and characterise the hidden SN populations in the nearest LIRGs. Such studies are needed to constrain the complete rates, properties and extinction distributions towards the CC SNe buried in such dusty galaxies. Eventually, these studies will hopefully provide a robust estimate for the numbers of SNe missed by optical searches both locally and at high-z. The resulting rate for the Ia supernovae is: 1.29$_{-0.57\,-0.28}^{+0.88\,+0.27}\times 10^{-4}$ yr$^{-1}$ Mpc$^{-3}$ h$^3_{70}$ at $\langle z\rangle =0.62$. This rate has been corrected for host galaxy extinction, but not for high redshift obscuration (a positive correction for this is included in the systematic error). Because of the quite high statistical errors and the lack of Ia SNe beyond $z=1$ we do not try to fit any DTDs to our Ia rate measurements. The comparison to the plotted models show that all these models are consistent with our rates. The high rate of Ia SNe at $z\sim 0.5$ measured by SVISS is in strong agreement with the results of \citet{2008ApJ...681..462D}, and thus in support of a Gaussian-like, fairly wide, DTD for Ia SNe. However, it should be noted that our measurement is also consistent with the distributions proposed by other authors. The type Ia SN rate measurements at $z\gtrsim 0.4$ differ by up to a factor of $\sim 2$. We believe that the cause of this large scatter is the large statistical and systematic errors that the SN surveys suffer from in this redshift range. It should also be noted that there are differences in the methods used to calculate the rates in the different surveys. It is therefore of utmost importance that the systematic errors are correctly estimated. \citet{2010ApJ...713...32S} find that models with a prompt Ia component are hard to reconcile with the rates measured at redshifts higher than one. Increased sample sizes at these high redshifts, or more studies of Ia host galaxies \citep{2005ApJ...634..210G,2006ApJ...648..868S} are needed to constrain the contribution from this channel. In Section~\ref{sec:debiasIa} we outline our motivation for not de-biasing the type Ia SN rates to account for extremely high extinction in U/LIRGs. It is very important to understand that the missing fractions of type Ia SNe situated in U/LIRGs similar to Arp~299 very likely depend on the DTDs. It is likely that also the missing fractions in other types of galaxies will depend on the DTD, although to a lesser extent. Using the type Ia SN rates to estimate the DTD thus inherently suffers from circularity to some extent. Assuming a zero missing fraction is a choice in itself -- the DTD implicit to such an assumption has a cut-off at some delay time $\tau \lesssim 200$ Myrs. If the DTD has no such cut-off, and a significant number of type Ia SNe have very short delay times, our assumption is faulty. In that case the resulting rates may be off by more than the adopted systematic error. The determination of supernova rates at high redshift is difficult. The SNe detected at high redshift will only be sampling the bright end of the SN luminosity function. This means that any global statistics estimated from such measurements will be sensitive to assumptions on the luminosity function made during the calculations. There are a number of additional assumptions going into the rate calculations that affect the rates differently. It is important to estimate the systematic errors these assumptions give rise to. With the exception of the misclassification error, that essentially scales with the sample size for samples smaller than ten given a misclassification ratio of 10\%, we have shown that the systematic errors are on the order of 50\% when using photometric redshifts and with the present uncertainties in template fractions and peak magnitudes. Furthermore, the assumptions made when correcting the rates for extinction/obscuration are shown to have a large effect on the systematics. Presently little is known about the number of SNe missed in LIRGs and ULIRGs, especially important at high redshift, which is evident in the large uncertainties on the de-bias factors. Given the low numbers of SNe for most high redshift surveys it is perhaps tempting to try and use the rates found in different surveys together when comparing to models (for the Ia SNe) and other sources (SFH for CC SNe). However, it's not straight-forward to do this. Different surveys estimate systematic errors in different ways, some include more sources and some less. If combined as given using co-added statistical and systematic errors, the risk is that greater weight is given to surveys where fewer systematic error estimates are included (given that the sample sizes are similar). We believe that the work presented in this paper shows the importance of including a variety of systematic effects to correctly estimate the uncertainties of supernova rates at high redshift. This will be even more important for future surveys with larger sample sizes and therefore lower statistical errors. | 12 | 6 | 1206.6897 |
1206 | 1206.6551_arXiv.txt | The background noise between 1 and 1.8 \um\ in ground-based instruments is dominated by atmospheric emission from hydroxyl molecules. We have built and commissioned a new instrument, GNOSIS, which suppresses 103 OH doublets between $1.47$ -- $1.7$\um\ by a factor of $\approx 1000$ with a resolving power of $\approx 10,000$. We present the first results from the commissioning of GNOSIS using the IRIS2 spectrograph at the Anglo-Australian Telescope. We present measurements of sensitivity, background and throughput. The combined throughput of the GNOSIS fore-optics, grating unit and relay optics is $\approx 36$ per cent, but this could be improved to $\approx 46$ per cent with a more optimal design. We measure strong suppression of the OH lines, confirming that OH suppression with fibre Bragg gratings will be a powerful technology for low resolution spectroscopy. The integrated OH suppressed background between 1.5 and 1.7 $\mu$m is reduced by a factor of 9 compared to a control spectrum using the same system without suppression. The potential of low resolution OH suppressed spectroscopy is illustrated with example observations of Seyfert galaxies and a low mass star. The GNOSIS background is dominated by detector dark current below 1.67 $\mu$m and by thermal emission above 1.67 $\mu$m. After subtracting these we detect an unidentified residual interline component of $\approx 860 \pm 210$ \bright, comparable to previous measurements. This component is equally bright in the suppressed and control spectra. We have investigated the possible source of the interline component, but were unable to discriminate between a possible instrumental artifact and intrinsic atmospheric emission. Resolving the source of this emission is crucial for the design of fully optimised OH suppression spectrographs. The next generation OH suppression spectrograph will be focussed on resolving the source of the interline component, taking advantage of better optimisation for a fibre Bragg grating feed incorporating refinements of design based on our findings from GNOSIS. We quantify the necessary improvements for an optimal OH suppressing fibre spectrograph design. | Observations at near-infrared wavelengths are severely hindered by the night-sky background. The night-sky surface brightness is $\approx 14.9$ AB mag arcsec$^{-2}$ in the H band, compared to $\approx 21.1$ AB mag arcsec$^{-2}$ in the V band. This results in high Poisson noise in any observation. Compounded with this, the night-sky brightness varies by factors of $\approx 10$ per cent on the timescale of minutes (\citealt{ram92}; \citealt{fre00}), with a further gradual dimming of about a factor of 2 throughout the night (\citealt{shi70}). This temporal variability results in a high systematic noise when performing sky subtraction, which is not trivial to remove (e.g.\ \citealt{dav07}; \citealt{sha10}). In the J and H bands the dominant sources of background are the Meinel bands of emission lines resulting from the rotational and vibrational de-excitation of OH molecules (\citealt{mei50}; \citealt{duf51}). We have reviewed the near-infrared background, and in particular the characteristics of the OH emission line spectrum in an earlier paper (\citealt{ell08}). We note that although they are very bright, the OH lines are intrinsically narrow (FWHM $\approx 3 \times 10^{-7}$\um), and between the OH lines the night-sky background should be very faint, possibly as low as $\approx 100$ \bright\ if the background is dominated by zodiacal scattered light, and is likely to be at least as faint as $\approx 600$ \bright\ as shown by $R=17,000$ spectroscopic observations made by \citet{mai93}. Thus, if the OH lines can be efficiently filtered from the night-sky spectrum, whilst maintaining good throughput between the lines, it should be possible to achieve a dark background in the NIR, allowing much deeper observations than possible hitherto. This paper presents the first results from an instrument designed to achieve exactly this efficient filtering of the OH skylines. The filtering is achieved using fibre Bragg gratings (FBGs). FBGs were originally developed for use in telecommunications, and required significant modification to be used for OH suppression (see \citealt{bland11} for a comprehensive treatment on the physics of OH suppression with FBGs). Firstly, it was necessary to significantly increase the number of notches in each grating (typically only one notch is used in devices used in telecommunications) and the wavelength range of the devices which were available at the time, a breakthrough which was enabled with the design of aperiodic FBGs (\citealt{bland04}). Secondly it was required to develop a multi-mode to single-mode fibre converter (\citealt{leo05,noo09}), since at the plate-scale of typical telescopes the narrow core of a single mode fibre has too small a field-of-view to collect an adequate amount of light from the seeing disc. Subsequent refinement (\citealt{bland08,bland11}) has resulted in the latest FBGs being able to suppress 103 notches over a wavelength range of 230~nm, which is achieved using two devices in series. \cite{ell08} presented the potential benefit for astronomy of OH suppression with FBGs and described the expected performance. \citet{ell08} point out that a key benefit of OH suppression with FBGs compared to other previously suggested methods of OH suppression (e.g.\ high dispersion masking, \citealt{mai00}) is that the OH light is removed before it enters the spectrograph and in a manner dependent only on wavelength. Hence the interline continuum is not contaminated with scattered OH light within the spectrograph which can otherwise dominate other sources of interline continuum. The testing of this prediction was a central goal of GNOSIS commissioning. The first on-sky tests of OH suppression with FBGs were performed at the Anglo-Australian Telescope in December 2008. This experiment consisted of two multimode fibres pointed directly at the sky through hole in the AAT dome wall. Both fibres fed a $1 \times 7$ photonic lantern (see \S~\ref{sec:inst}), one of which had FBGs inserted, and the other did not, to serve as a control. The results of these tests are described in \citet{bland09,bland11}, and demonstrated the clean suppression of the OH lines. However, since each fibre accepted light from a $\approx 12$ degree patch of sky these tests could not perform observations of individual sources, nor could they measure the interline continuum since every observation included light from many stars and other sources. We have designed, built and commissioned the first instrument to use FBGs for OH suppression. GNOSIS is an OH suppression unit designed to feed the IRIS2 infrared imaging spectrograph at the AAT (\citealt{tin04}). The name is an acronym for Gemini Near-infrared OH Suppression IFU System, revealing the intention to install a similar system on the GNIRS spectrograph at Gemini North (\citealt{eli98,eli06,eli06b}) at a future date. A summary description of the instrument is given in section~\ref{sec:inst} below. In this paper we concentrate on the performance and results of the OH suppression. We describe the observations in section~\ref{sec:obs} and we present the results in section~\ref{sec:results} giving details of throughput, sensitivity and background. The near-infrared background will be examined in detail in a future paper (Trinh et al., in prep.), examining the various components of the background and their dependence on moonlight, ecliptic latitude, airmass, Galactic latitude etc.; a brief description of the main components will be given here. In section~\ref{sec:sciobs} we present observations of two Seyfert galaxies and an L/T dwarf illustrative of OH suppression. In section~\ref{sec:disc} we discuss our results in the context of the benefit of OH suppression to NIR spectroscopy and in the context of lessons for future OH suppressed instruments. | \label{sec:disc} We have demonstrated the first instrument to employ OH suppression with fibre Bragg gratings. To test the performance of this new technology we made five key measurements: (i) the overall, and component, instrument throughput, (ii) the instrument sensitivity, (iii) the level of OH suppression, (iv) the interline component, (iv) illustrative observations of Seyfert galaxies. We now discuss the results of these measurements in the context of the continuing development of OH suppressed NIR spectrographs. (i) The total throughput is $\approx 4$ per cent but the throughput of IRIS2 is only $\approx 12$ per cent while that of GNOSIS itself is $\approx 36$ per cent. The latter can be greatly improved by the following measures. First on a specialised OH suppression spectrograph the need for relay optics could be eliminated by using a fibre vacuum feed through. This would remove at least 4 optical surfaces, greatly improve alignment losses and have the additional benefit of lowering the thermal background. Secondly, an even larger improvement could be made in the photonic lanterns. Recent laboratory tests imply that the efficiency of the photonic lanterns is a function of the input f ratio, and that by feeding the lanterns at a slower speed we could increase the efficiency from 50 per cent to $\approx 70$ per cent. These tests are preliminary, but indicate that the coupling efficiency of the lanterns may depend on the mode being coupled. (ii) The sensitivity of GNOSIS+IRIS2 is about that of IRIS2 when used in its standard spectrograph mode with a slit.. That is, in the particular implementation used for GNOSIS, the benefits of OH suppression are offset by lower throughput, higher detector background, higher thermal background and lower observing efficiency. The low throughput has been discussed above and in section~\ref{sec:through}. The high detector background comes from the fact that a small field of view of 1.2 arcsec is being divided into 7 elements, and the spectra from each of these is being sampled by $\approx 2$ pixels, thus when extracting the spectra for each spectral pixel there is dark current and read noise from $\approx 14$ spatial pixels, compared to $\approx 2$ spatial pixels for the same sky area using IRIS2 (recall that GNOSIS feeds IRIS2 at a slower f-ratio than for long-slit IRIS2 observations). Since the background is dominated by detector dark current this is a significant loss of sensitivity. The high thermal background results almost entirely from the slit block, and is a significant source of noise at $\lambda \gsim 1.65 \mu$m. GNOSIS observations are also less efficient than IRIS2 observations in that half the time must be spent on sky, whereas with IRIS2 the object can be nodded up and down the slit. Note that these problems are not intrinsic to OH suppression in general, and it is easy to envisage improved instruments that circumvent all these issues. Using photonic lanterns with a larger number of cores will allow the use of larger diameter fibres and thus increase the field of view per fibre, thereby lessening the effects of detector background. This will require faster spectrograph optics of around $\approx f/2.8$. The low observing efficiency is simply a matter of cost. If two IFUs can be afforded then cross beam switching of observations will ensure that all observing time is spent on target. Of course there are also many improvements that should be made to the spectrograph itself to improve the overall efficiency. For example newer detectors with lower dark current and higher quantum efficiency, VPH gratings instead of grisms, and fixed format optics will all improve the throughput. These improvements would of course be true of non-OH suppressed spectrographs as well. (iii) The suppression of the OH lines was a success. The sky spectra and the object spectra show that the OH lines can be cleanly removed whilst maintaining good throughput (in the FBGs themselves) between the lines. The level of suppression is high with $\approx 78$ per cent of the targeted lines suppressed at the required level. Of the lines which do not meet the required level of suppression 37 per cent are due to residuals from the Q1(3.5) and Q1(4.5) OH lines (\S~\ref{sec:ohsupp}). Similarly there is an unsuppressed O$_{2}$ line. (iv) The expected reduction of the interline component was not observed. This seems to indicate either an unaccounted source of interline emission, inaccuracy of our OH line models or unaccounted-for systematic errors. Systematic errors are a real possibility since we are operating in a very low-count regime of less that 1 e$^{-}$ per pixel per minute, and we are dominated by detector noise. We have observed (and corrected for) systematic deviations from detector linearity as discussed in section~\ref{sec:cube}, and have noted that there may still be other effects such as reciprocity failure (\citealt{bie11}) which are very difficult to characterise. We have already noted that specific transitions in the OH energy levels were not properly suppressed, which seems to be due to a larger energy gap between the $\Lambda$ doublets than predicted in our models. It is possible that other such errors may exist in fainter lines, or that other branches of lines are brighter than supposed, but we are unable to measure this with the current sensitivity. Unsuppressed faint OH lines could masquerade as continuum due to the spectrograph scattering. On the other hand there could be emission from another source. A thorough search of the emission lines from 33 different molecules in the HITRAN2008 database cannot account for the structure and the faint lines we observe in the interline component except for a couple of individual features from O$_{2}$ and CH$_{4}$. Also there is no dependence of the interline component on airmass, as would be expected for an atmospheric source. However, there always remains the possibility that there is another overlooked source or that or relative intensities of the molecular emission needs improving. To try to distinguish between these two possibilities we looked at the correlation between the interline component and the OH line strength and found that they were only weakly correlated. If true, this suggests that there is indeed an unknown source of interline emission, since otherwise the correlation would be much stronger. However, this interpretation should be treated with much caution due to the large measurement errors and possibly significant systematic errors on the interline measurements. The absolute level of the interline emission is $860 \pm 210$ \bright\ which is slightly higher than that measured by \citet{mai93}, (590 \bright) and slightly less than that measured by \citet{cub00} (1200 \bright). We again point out that there may be significant systematic error on our measurement due to the low count rates measured. However, the similarity to previous measurements could indicate that we have indeed reached a floor in the interline component which is significantly higher than the zodiacal scattered light background (see \citealt{ell08}). If the interline emission is the result of unsuppressed OH lines due to inaccuracies in our OH line models, this may be fixed with better models or with wider FBG notches to account such uncertainties. However, as things stand, the lack of a reduction in the interline component suggests that OH suppression will find its niche in low resolution ($R \lsim 3000$) observations which would normally be insufficient to `resolve out' the OH lines. Indeed OH suppression could be used at low resolutions, $R \sim 500$, at which conventional observations would be impractical. (v) The suggested benefit of OH suppression at low resolution is supported by our observations of Seyfert galaxies (\S~\ref{sec:sciobs}). We find that in regions where the OH line density is high, as for NGC~7674, the low resolution of IRIS2 makes it difficult to identify the [FeII] emission in the control spectrum, but that the line is obvious in the suppressed spectrum. Similarly when working between the lines as for NGC~7714, then the benefits of OH suppression are much less apparent since in this case we are limited by the detector dark current in both cases. We conclude by noting that GNOSIS is the first step toward an OH suppression optimised spectrograph, and not a fully fledged facility instrument. The overall performance of the OH suppression unit itself is good, but can be substantially improved. However there remains uncertainty about the origin of the interline emission, which could higher than predicted either due to some unknown source of emission or because of unaccounted-for systematics in the detector. The next step will be to design and build a spectrograph optimised for a FBG feed, which will avoid the disadvantages that result from retrofitting an OH suppression system to an existing spectrograph. Furthermore, in doing so we can capitalise on the lessons learnt with GNOSIS in terms of detector background, optimal spectral resolution and design. We present simulations of such a system in Figure~\ref{fig:sims}. This simulation is based on the method presented in \citet{ell08} with modifications to reflect our experience with GNOSIS. These modifications include an updated OH line spectrum which is an amalgamation of the observations of \citet{abr94} with the models of \citet{rous00}, giving precedence to \citet{abr94} where lines occur in both lists. The thermal background modelling has also been updated to consider the emission and the numerical aperture of each component of the optical train separately. We first use the models to simulate the GNOSIS spectrum and compare these to the observations. We next model PRAXIS, a concept for an optimised OH suppression system. Specifically we assume that we have a corrected FBG design to account for the residual OH lines, that the fibre slit is cooled via a vacuum feed-through, that we have improved spectrograph throughput through the use of VPH gratings, and lower intrinsic detector noise through the use of a Hawaii-2RG 1.7$\mu$m cut-off detector, and better spatial sampling with 2 pixels per PSF FWHM. | 12 | 6 | 1206.6551 |
1206 | 1206.5308_arXiv.txt | We study the properties of two bars formed in fully cosmological hydrodynamical simulations of the formation of Milky Way-mass galaxies. In one case, the bar formed in a system with disc, bulge and halo components and is relatively strong and long, as could be expected for a system where the spheroid strongly influences the evolution. The second bar is less strong, shorter, and formed in a galaxy with no significant bulge component. We study the strength and length of the bars, the stellar density profiles along and across the bars and the velocity fields in the bar region. We compare them with the results of dynamical (idealised) simulations and with observations, and find, in general, a good agreement, although we detect some important differences as well. Our results show that more or less realistic bars can form naturally in a $\Lambda$CDM cosmology, and open up the possibility to study the bar formation process in a more consistent way than previously done, since the host galaxies grow, accrete matter and significantly evolve during the formation and evolution of the bar. | Bars are present in roughly two thirds of all galactic discs \citep{Eskridge00,Barazza08} and can drive the secular evolution of their host galaxies. They have thus been the subject of a number of studies based on dynamical simulations, i.e. simulations aimed towards an understanding of the main relevant dynamical mechanisms (e.g. \citealt{Combes90,Debattista2000,AM02, A03, MartinezValpuesta06}). These use idealised initial conditions, made specifically for the question under study, and set aside all other effects, at least until an understanding of the basic mechanism is achieved. The initial conditions correspond to a fully developed disc plus halo system which is as near equilibrium as possible, so that the bar formation can be studied uninfluenced by other instabilities. Furthermore, interactions with other galaxies, major or minor, as well as inflow from the environment has been seldom taken into account (see, however, \citealt{Curir06}), while the halo is usually assumed to be initially spherical. Moreover, the gas is generally neglected or modelled without taking into account star formation, feedback or cooling. In this letter, we take a different approach, i.e. we study bar formation in the context of the $\Lambda$CDM cosmology. In this way, we include important external effects, such as accretion and interactions and, more important, the bars do not wait for the galaxy to be fully formed to start their own formation. The disc and even the dark matter halo keep growing while the bar forms. Also, our galaxies do not have preset halo-to-disc mass ratios, velocity dispersions, or other properties, and all their properties directly result from the simulations. This of course implies that we will not be able to address the same questions as the dynamical simulations. For example, we can not examine how a given property of the host galaxy, e.g. the halo radial density profile, will influence the formation and evolution of the bars. On the other side, we will include in our simulations more physics that in any other single bar formation simulation. Our main goal is to investigate if bars can naturally form in $\Lambda$CDM, and compare their properties with those of dynamic simulations and with observations. The layout of the paper is as follows. In Section~\ref{sec:sims} we describe the simulations, in Section~\ref{sec:results} we present and discuss our results, and we conclude in Section~\ref{sec:conclusions}. \begin{figure} {\includegraphics[width=36mm]{figures/stars_xy_1_128.eps}\includegraphics[width=36.mm]{figures/stars_xy_0_128.eps}} {\includegraphics[width=36mm]{figures/stars_xz_1_128.eps}\includegraphics[width=36mm]{figures/stars_xz_0_128.eps}\includegraphics[width=36mm,angle=90]{figures/ColorBar_isodensities.eps}} {\includegraphics[width=36mm]{figures/stars_yz_1_128.eps}\includegraphics[width=36mm]{figures/stars_yz_0_128.eps}} \caption{Projected surface stellar density for Aq-C (left-hand panels) and Aq-G (right-hand panels). From top to bottom, the figures show the face-on ($XY$), side-on ($XZ$) and end-on ($YZ$) views. The black lines correspond to isodensity contours equally-spaced in log$\Sigma$, and the white circles indicate the bar length obtained with three different methods (see Section~\ref{sec:length}). } \label{fig:isodensities_z0} \end{figure} \vspace{-0.2cm} | \label{sec:conclusions} We presented the first study of bars formed in fully cosmological, hydrodynamical simulations of Milky Way-sized haloes. % The simulations include star formation, metal-dependent cooling, feedback from supernova and a UV background field. In particular, we investigated the morphology, strength and length of the bars, their projected density profiles and kinematical properties. The strongest of our two bars formed in a bulge-disc-halo system (Aq-C), while the weakest in a galaxy with a disc but no significant bulge (Aq-G). Compared to dynamical simulations, the strongest bar is similar to those found in systems where a considerable amount of angular momentum is exchanged between the bar and the halo. In contrast, the weakest bar is more reminiscent of dynamical simulations where considerably less angular momentum has been redistributed. The bar strength difference between our two simulated bars could be due to the effect of the bulge on the angular momentum exchange (which is negligible in Aq-G), as found in pure N-body dynamical simulations of isolated discs (AM02, A03), but the effect of interactions and/or of the gas component could also be decisive. The Aq-C (strongest) bar is very thin and rectangular-like, while that of Aq-G (weakest bar) is very fat, and also rectangular-like. Rectangular shapes are observed in strongly barred galaxies, so Aq-G is, in this respect, like an hybrid between strong and weak bars. Another difference is that the bulge of Aq-C has a rather peculiar shape which has so far not been observed. The peculiar shape is presumably a result of the cosmological formation of our galaxies, where mergers and/or interactions are common at all times. The lengths of our bars are in relatively good agreement with observations of G11, although in the long tail of the distribution. The strongest bar, formed in our most massive galaxy, is also the longest. On the other hand, the ratios between bar length and disc scale-lenght are at the low end of the observed distribution, but still compatible with it. The density profiles are similar to those found in dynamical simulations; however, we detect some differences particularly in Aq-C and in the very central regions, probably because we include stars from all components in our analysis, unlike in dynamical simulations. Finally, the kinematic properties of our bars are similar to those observed and to those found in dynamic simulations. The fact that the bars were obtained in a cosmological setting, and that their properties agree relatively well with known properties of bars, is not minor. Our galaxies grow significantly from $z=3$ to $z=1$, where the bars are first detected in the simulations. Moreover, external (accretion, interactions and mergers) and internal (cooling, star formation and feedback) effects strongly affect the galaxies during their evolution. For these reasons, our results are quite satisfactory and very encouraging, and should incite more work on the study of bar formation in a cosmological context. \vspace{-0.3cm} | 12 | 6 | 1206.5308 |
1206 | 1206.2833_arXiv.txt | We use numerical models, supported by our laboratory data, to predict the dust densities of ejecta outflux at any altitude within the Hill spheres of Europa and Ganymede. The ejecta are created by micrometeoroid bombardment and five different dust populations are investigated as sources of dust around the moons. The impacting dust flux (influx) causes the ejection of a certain amount of surface material (outflux). The outflux populates the space around the moons, where a part of the ejecta escapes and the rest falls back to the surface. These models were validated against existing Galileo DDS (Dust Detector System) data collected during Europa and Ganymede flybys. Uncertainties of the input parameters and their effects on the model outcome are also included. The results of this model are important for future missions to Europa and Ganymede, such as JUICE (JUpiter ICy moon Explorer), recently selected as ESA's next large space mission to be launched in 2022. | Micrometeoroids (solid micron-sized dust particles) are a common constituent of the Solar System. They can easily reach the surfaces of atmosphereless bodies and, upon impact, cause the ejection of surface material into the surrounding space. The ejected dust fragments populate the space around the host bodies, where a part of the ejecta escapes and the rest eventually falls back to the surface. This process can be numerically modelled. In this paper we investigate dust around Europa and Ganymede created by micrometeoroid bombardment. \\* Our study is built upon the models developed by \citet{Krivov2003} and \citet{Kruger2003} supported by our impact experiments \citep{Miljkovic2011a} and hydrocode impact modelling to reduce the number of variables in the model. The impact experiments were made using the light gas gun (LGG) at the Open University's Hypervelocity Impact (HVI) laboratory \citep{Miljkovic2011a, Patel2010, McDonnell2006, Taylor2006}. A series of high velocity impacts were made at $2\,km\,s^{-1}$ using 1 mm diameter stainless steel balls as projectiles into pure water ice and sulphate hydrated mineral targets. The ejecta size \citep{Miljkovic2011a} and velocity distributions were measured and subsequently modelled using ANSYS AUTODYN finite element hydro-dynamic shock physics code \citep{Miljkovic2010, Pierazzo2010}. \\* The dust cloud model presented in this paper characterizes the dust environment (size, density, flux and velocity distribution of such dust) around Europa or Ganymede, predicts the dust densities of the ejecta outflux at any altitude within Europa's and Ganymede's Hill spheres (radius of $13300\,km$ or $8.5\,R_{E}$, where $R_{E}=1565\,km$ for Europa and $32000\,km$ or $12\,R_{G}$, $R_{G}=2631\,km$ for Ganymede) and can evaluate the dust density at any altitude as a function of the size distribution of the dust. We choose a discrete selection of altitudes and dust masses to give a representative set of results. \\* Our results are important for future space missions to Jupiter System that carry a dust detector onboard. This study can be further applied to estimate the dust counts into a dust detector in orbit around Europa and/or Ganymede. A dust detector has been proposed as part of a payload for a space mission to Europa and Ganymede. Initially named Laplace in 2007, it was renamed the Europa-Jupiter System Mission after ESA and NASA joined proposals in 2009 \citep{Blanc2009} for a major mission to Jupiter System. In 2011, EJSM was reformed again into an ESA-led mission to Ganymede with flybys to Europa and Calisto, named JUICE (JUpiter ICy moon Explorer) \citep{Dougherty2011}. JUICE has recently been selected as ESA's next large mission for launch in 2022. If a dust detector is included in the payload, an in-situ analysis of the dust that surrounds Europa and Ganymede will be possible, which could provide information about the surface, as its composition should be "written" in the detected dust \citep{Miljkovic2011b}. \\* The proposed dust detector should not only be capable of determining the density of dust in the cloud, but may provide chemical analysis of captured dust \citep{Miljkovic2008}, as was the case with the analysis of the Jovian Dust Stream particles by Cassini's CDA (Cosmic Dust Analyser) \citep{Postberg2006}. Chemical abundance maps of Europa and Ganymede even at low spatial mapping resolution show that water ice is non-uniformly distributed over the surfaces of Europa and Ganymede. The Galileo NIMS (Near Infra-red Mapping Spectrometer) spectra identified these impurities as hydrated minerals, sulphates and possibly hydrocarbons \citep{McCord1998}. As the surface material is ejected by micrometeoroid bombardment, it can be expected that the dust particles around Europa will be composed of water ice, sulphur salts and their decomposition products, including any potential organic compounds \citep{Miljkovic2011b}. It should be emphasized that a dust detector has never visited Europa or Ganymede at such a close orbital distance or spent a longer time than in flybys. No chemical analysis of the dust in Jupiter System has been made yet, apart from Cassini's Cosmic Dust Analyser (CDA) measurements of the Io stream dust at more than $1\,AU$ away from Jupiter \citep{Postberg2006}. \subsection{Comparison of the present dust cloud model with previous models.} In previous work, \citet{Krivov2003} developed a spherically symmetric case for an atmosphereless body with applications to Ganymede and the Saturnian satellites to help the interpretation of Cassini's CDA (Cosmic Dust Analyser) measurements; \citet{Kruger2000} developed a dust cloud model around Ganymede to explain the Galileo DDS measurements; \citet{Sremcevic2003} investigated the asymmetry of the dust clouds around Galilean and Saturnian satellites; \citet{Sremcevic2005} compared the Galileo data with their models for Europa, Ganymede and Callisto, whereas \citet{Kruger2003} investigated the dust clouds around all four Galilean satellites. \\* The differences between the dust model presented here and that of \citet{Kruger2003} are the following: their slope of the cumulative ejecta mass distribution was approximated, whereas in our model our slope was derived directly from the impact experiments published in \citet{Miljkovic2011a}. \citet{Kruger2003} reported that dust have mean mass of $10^{-11}$g around Europa and $10^{-13}$g around Ganymede, whereas in our model the ejecta fragment size distributions were calculated, and were in range between $10^{-15}$g-$10^{-7}$g. Ejecta speed distributions in \citet{Kruger2003} were represented as the distribution of ejecta material having speeds higher than a certain speed, that is dependent on the minimum ejecta speed of fragments and the power-law slope of the distribution, which were unknown variables fitted to match the Galileo data, whereas in our model, a size-velocity relation was applied to all the ejected fragments in order to have more precise outflux and spatial densities of ejected dust. We primarily focus on the short-lived, bound ejecta at a distance from the surface at which a spacecraft may orbit. Therefore, any asymmetry effects in the spatial density of ejected fragments caused by Europa's orbital motion can be excluded (such were considered by \citet{Sremcevic2003}) as well as the charging of dust fragments. | In this paper, we present an update model for determining the size and spatial density of dust around Europa and Ganymede at any altitude within the moons' Hill sphere, with particular emphasis on the role of micrometeoroid impacts ejecting surface material. The model computes influx distributions from different sources and addresses each ejected fragment individually according to its size and corresponding speed, allowing the complete profile of the dust around Europa and Ganymede to be defined through calculation of the size, velocity and spatial distribution of the complete outflux. \\* An important feature of this dust model is the ability to predict the dust around Europa for distances closer to the surface than any of the spacecraft flybys made so far as well as for the dust populations below the Galileo DDS detection threshold. We have verified these results by comparing our results with data collected by the Galileo DDS. It is found that the dust spatial densities close to the surface of Europa and Ganymede should be much higher than a simple extrapolation from the actual Galileo data. There should also be much more dust at sizes below the Galileo DDS detection threshold. \\* This model can be applied to any other atmosphereless body in the Solar System, but above all, it is of great importance for future orbiter missions around Europa and Ganymede, such as recently selected JUICE mission, ESA's next large space mission, to Jupiter System. This study provides scientific support for a dust detector/analyser payload for JUICE or any other future space mission to Europa and/or Ganymede. It also complements surface studies traditionally performed using remote sensing instruments. | 12 | 6 | 1206.2833 |
1206 | 1206.0007_arXiv.txt | The first population III stars are predicted to form in minihalos at $z \approx 10$--30. The {\it James Webb Space Telescope (JWST)}, tentatively scheduled for launch in 2018, will probably be able to detect some of the first galaxies, but whether it will also be able to detect the first stars remains more doubtful. Here, we explore the prospects of detecting an isolated population III star or a small cluster of population III stars down to $z=2$ in either lensed or unlensed fields. Our calculations are based on realistic stellar atmospheres and take into account the potential flux contribution from the surrounding H\,\textsc{ii} region. We find that unlensed population III stars are beyond the reach of {\it JWST}, and that even lensed population III stars will be extremely difficult to detect. However, the main problem with the latter approach is not necessarily that the lensed stars are too faint, but that their surface number densities are too low. To detect even one 60 M$_{\odot}$ population III star when pointing {\it JWST} through the galaxy cluster MACS J0717.5+3745, the lensing cluster with the largest Einstein radius detected so far, the cosmic star formation rate of population III stars would need to be approximately an order of magnitude higher than predicted by the most optimistic current models. | \label{introduction} Both theoretical arguments and numerical simulations \citep[e.g.][]{2004ARA&A..42...79B}, strongly support the notion that population III stars were very massive, significantly more so than the population I and population II stars that formed later on. When stars form in metal-enriched gas, the Jeans mass is lower, which leads to fragmentation and thus lower masses. Lacking metals, a chemically unenriched cloud does not fragment in the same way and the characteristic stellar mass may therefore be higher. It has been argued that two different classes may have existed: population III.1 and population III.2. Population III.1 stars formed in dark matter minihalos of mass 10$^5$--10$^6$ M$_{\odot}$ at $z\approx$ 10--30. Early research on population III star formation \citep{2007ApJ...654...66O, yoshidaomukaihernquist2008, 2009Natur.459...49B}, indicated that only one star, with a very high average mass of $\sim$ 100 M$_{\odot}$ was produced in every minihalo. This was the result of UV radiation (in the Lyman-Werner band) produced by the first massive star destroying the molecular hydrogen in the parent cloud, preventing further cooling and star formation. Later research on the same topic \citep[e.g][]{2009Sci...323..754K, 2009Sci...325..601T, 2010MNRAS.403...45S, 2011ApJ...731L..38P, 2011ApJ...727..110C, 2011ApJ...737...75G, 2012MNRAS.422..290S} indicates a more complex scenario with more fragmentation, resulting in a binary system or small cluster of stars, with the most massive star reaching $\approx$~50~M$_{\odot}$. \citet{2008ASPC..393..275S} also suggest that cosmic rays from the supernovae resulting from the first massive population III stars could significantly influence subsequent star formation in minihalos, possibly yielding a characteristic stellar mass of $\sim$~10~M$_{\odot}$. It is also plausible that protostellar feedback will halt the mass accretion, thereby limiting the stellar mass to a few times 10~M$_{\odot}$ \citep{2011Sci...334.1250H}. There are also models indicating late, $2 < z < 7$, population III star formation \citep{2007MNRAS.382..945T, 2010MNRAS.404.1425J}. These late population III stars form in regions that are pristine due to highly inhomohegeneous metal enrichment. These low-metallicity regions are usually located in the vicinity of more massive ($\sim 10^8-10^{11}$~M$_{\odot}$) dark matter halos containing metal enriched galaxies. In \citet{2011MNRAS.411.2336I}, the authors examine LAEs at $z=3.1$ with signatures they argue are hard to interpret as other than objects containing population III stars, (but see \citet{2012MNRAS.tmp..127F} for a different view). Population III.2 stars are thought to form mostly due to cooling by adiabatic expansion or HD cooling, usually at a lower redshift than population III.1 stars, and are thought to be building blocks for early galaxies. Since we focus on pre-galactic population III star formation we will not investigate these any further. \citet{2001ApJ...552..464B} have investigated the spectral energy distribution of a primordial star. \citet{2006SSRv..123..485G} and \citet{2009MNRAS.399..639G} used these results to show that isolated population III stars are too faint for detection with the {\it JWST}. However, the flux boost due to gravitational lensing has not previously been considered in detail, except for a preliminary version of the study we present here \citep{2010crf..work...26R}. Apart from direct detection through gravitational lensing, there are other ways to probe the properties of population III stars. The cumulative indirect signatures in the global radio background produced by bremsstrahlung or the 21-cm hydrogen emission also represent interesting options. Bremsstrahlung is free-free emission originating in the H\,\textsc{ii} region in its active and relic states. When the population III star has died the almost fully ionized H\,\textsc{ii} region around it begins to recombine. This makes it a very bright source of hydrogen 21-cm radiation \citep{2009MNRAS.395..777T}. \citet{2009MNRAS.399..639G} find that bremsstrahlung is very unlikely to be detected in the near future but the cumulative 21-cm radiation might be strong enough to be within reach of the upcoming Square Kilometer Array. Direct observation of pair instability supernovae (PISN) could also be used to infer the existence of population III stars. PISNs are the predicted final fate of non-rotating stars in the mass range 140--260~M$_{\odot}$ \citep{1967PhRvL..18..379B, 1968Ap&SS...2...96F, 2012arXiv1211.4979W, 2012MNRAS.tmp..244D}, and are possible for lower mass stars if rotating \citep{2012ApJ...760..154C}. The initial mass function (IMF) describes the distribution of mass for a new star. Since PISN occur in a certain mass range observations of those could probe the IMF. Using a semi-analytic halo mass function approach, \citet{2012ApJ...755...72H} find that no more than $\sim$~0.2 population III PISN should be visible per {\it JWST} field of view at any one time, depending on the adopted model for stellar feedback. Rapidly rotating population III stars may also produce collapsar gamma-ray bursts, or GRBs \citep{2011A&A...533A..32D}. These are formed by metal poor, rotating stars more massive than $\sim$~40~M$_{\odot}$. This could be a highly energetic probe of population III stars, and of the high-mass tail of their IMF. Population III stars could also leave detectable imprints \citep{2006astro.ph.10943K} in the cosmic infrared background (CIB). The CIB is a repository of emissions from different sources at different redshifts, all redshifted into the infrared. The imprints from population III stars are already constrained by detection of emission at energies above 10~GeV \citep{2012MNRAS.420..800G}, as interaction between the radiation of population III stars and gamma rays create a gamma-ray optical depth. As we will argue, a flux boost due to gravitational lensing by a foreground galaxy cluster will be necessary to directly detect population III stars. By adopting the properties of one of the more promising lensing clusters, we will calculate the SFR required to achieve $\sim$~1 detectable population III star per survey field. | \label{sec:discussion} Of the many scenarios explored in Section~\ref{sec:JWSTbroadbandfluxes}, the only one that would make isolated population III stars potentially detectable with {\it JWST} is: \begin{itemize} \item Very massive ($\gtrsim 300$~M$_{\odot}$) population III stars. \item The stars are formed in the limiting case of either no L-W feedback \citep{2009ApJ...694..879T} or in the prolonged population III star formation case predicted by \citet{2012ApJ...745...50W}. \item The Lyman-alpha escape fraction would need to be very high ($f_\mathrm{Ly\alpha} \approx 0.5$). \item The gravitational lens contains sufficiently large regions with very high magnifications ($\gtrsim$~1000). \item The escape of Lyman continuum (i.e. hydrogen-ionizing) photons from the nebula surrounding the population III stars is very small. \end{itemize} Hence, the prospects for detecting isolated population III stars in the foreseeable future appear bleak. However, contrary to previous claims \citep{2006SSRv..123..485G, 2009MNRAS.399..639G}, the problem is not necessarily that population III stars are too faint for detection, as the magnification of a suitably chosen lensing cluster may lift them above the {\it JWST} detection threshold. The main obstacle is instead that the surface number densities for sufficiently massive population III stars to be found in a sufficiently lensed field is likely to be too low. Even so, given the scenario that we have a detection of a population III star we still have to identify it as such. One plausible mean of doing this suggested in the literature \citep{2001ApJ...553...73O, 2001ApJ...550L...1T} is to identify the He\,\textsc{ii}~1640~\AA{} line. Population III stars are predicted to have a significantly harder spectrum than even low metallicity stars. This means significantly more He are ionized in relation to H ionized. The result is that the ratio of the He\,\textsc{ii}~1640~\AA{} line to a non-resonant hydrogen line such as H$\alpha$ should be significantly larger for population III stars compared to that from even low metallicity stars. It has also been argued that supersonic streaming velocities between dark matter and gas could delay population III star formation, \citep{2011ApJ...736..147G, 2011MNRAS.412L..40M}. The delay is substantial, $\Delta z \sim 4$, but the large shift in mass also means the number density of population III star-forming minihalos is reduced by up to an order of magnitude. There is criticism of the importance of the effect \citep{2011ApJ...730L...1S}, but if significant, it would of course impact the results presented here. The combined effect of fewer stars at lower redshift needs careful examination in future work. As mentioned in Section~\ref{sec:modellingspectrapopulationIIIstars}, the low densities expected in the H\,\textsc{ii} regions could make time dependence important, in particular when the time scales of ionization/recombination and cooling/heating become comparable to the lifetimes of the stars. We have tested this using two scenarios for the surrounding H\,\textsc{ii} regions, one with 1 and another with 100 hydrogen atoms~cm$^{-3}$. We set the He/H number ratio to 0.09, and the abundances of other elements to zero. We show the relative luminosity of Ly$\alpha$ for the two densities and for both the 60~M$_{\odot}$ and 300~M$_{\odot}$ stars in Figure~\ref{fig:lyatimedependence}. We have assumed that the stars suddenly switch on and shine for about 3.4~Myrs with a constant luminosity and spectrum. For $n_\mathrm{H} = 100$~cm$^{-3}$, the ionization quickly goes toward the steady-state solution, whereas for $n_\mathrm{H} = 1$~cm$^{-3}$ it takes $\sim 0.4$~Myrs to reach the steady-state level. However, even after this level has been reached, there is still some slow photoionization going on in the outer parts of the H\,\textsc{ii} region. The temperature in this region also overshoots somewhat compared to the steady-state solution before steady state has been attained, which produces slightly enhanced collisional excitation of Ly$\alpha$. This effect is larger for the 300~M$_{\odot}$ star with its more extended region of partially ionized hydrogen outside the Str\"omgren radius. The net effect of time dependent versus steady-state models is, as can be seen from Figure~\ref{fig:lyatimedependence}, not more than $\sim 10-15$~\%. For densities below $n_\mathrm{H} = 1$~cm$^{-3}$, the deviations from steady-state will obviously become larger. However, the H\,\textsc{ii} region is expected to be evolving from high to low density with time \citep[e.g.][]{2004ApJ...610...14W}. This probably makes the early delay of Ly$\alpha$ in Figure~\ref{fig:lyatimedependence} exaggerated as, at this time, the H\,\textsc{ii} region is dense. At low densities the emission from the H\,\textsc{ii} region will fall off on a time scale comparable to the lifetime of the star. For such low densities, the H\,\textsc{ii} region may however no longer even be ionization bounded. The uncertainty in our models due to time dependence is in any case much smaller than due to other assumptions. \begin{figure} \begin{center} \includegraphics[width = 10 cm]{TimeDependentRelative.eps} \caption{The flux in the Lyman alpha line from a $n_\mathrm{H} = 1$~cm$^{-3}$ nebula divided by the flux in the Lyman alpha line from a $n_\mathrm{H} = 100$~cm$^{-3}$ nebula, as a function of time. We have assumed that the stars suddenly switch on and shine for about 3.4~Myrs with a constant luminosity and spectrum. For the reference nebula with $n_\mathrm{H} = 100$~cm$^{-3}$ the ionization quickly goes toward the steady-state solution. The blue line represents a 60~M$_{\odot}$ star, whereas the purple line represents a 300~M$_{\odot}$ star. For $n_\mathrm{H} = 1$~cm$^{-3}$ it takes $\sim 0.4$~Myrs to reach the steady-state level. A 60~M$_{\odot}$ star's nebula subsequently emits $\sim$~5~\% more flux than in steady state and a 300~M$_{\odot}$ star's nebula up to $\sim$~15~\% more flux than in steady state. The upturn seen for both stellar models after 3.4~Myrs is an effect of time dependence in the low density (1~cm$^{-3}$) case. For both densities, the Ly-alpha emission starts to decay when the star is switched off, but the decay is about two orders of magnitude faster for 100~cm$^{-3}$ than for 1~cm$^{-3}$.} \label{fig:lyatimedependence} \end{center} \end{figure} We have calculated the required magnification for detection of population III stars using realistic models for the atmosphere and the H\,\textsc{ii} region. However, our results are admittedly based on a number of simplifications, most of them made in the direction that would improve the prospects of detection. The spherically symmetric H\,\textsc{ii} region that we have assumed likely overestimates the flux, as there probably will be low-density channels in the nebula through which ionizing radiation would be able to escape into the intergalactic medium. We have also assumed a spatially homogeneous distribution of minihalos. In reality clustering could impact the detection probability in a negative way. All in all, even though not impossible, the prospect of detecting population III stars (isolated, or in small clusters) with {\it JWST} appears bleak at best. Larger numbers of population III stars could potentially form in chemically pristine (10$^7$--10$^8$ M$_{\odot}$) atomic-cooling halos at $z<15$ \citep[e.g.][]{2012MNRAS.426.1159S}. These ``population III galaxies" would by comparison be much easier to detect and also identify \citep{2009MNRAS.399...37J, 2011ApJ...731...54P, 2011ApJ...740...13Z}. The detectability of such objects hinge more on the total star forming mass than on stellar masses and nebular properties, resulting in less extreme lensing magnification requirements \citep{2012arXiv1204.0517Z}. The merits of investigating the prospects for detecting population III stars in minihalos is that they possibly fill an important role in the development of the universe, providing the first metal-enrichment prior to the first galaxy formations \citep{2010ApJ...716..510G}. More speculative and exotic theories involving population III stars or derivatives of them also exist, potentially improving detection prospects. Fast accretion combined with rapid rotation could potentially form super massive population III stars \citep{2012ApJ...756...93H}. These would be extremely massive, up to 10$^5$--10$^6$~M$_{\odot}$, and have lower temperature, $\sim$~5000~K. A black hole could also form by collapse in the center of the star powering it through emission by accretion, a so-called quasi-star \citep{2010MNRAS.402..673B}. Some significant fraction of population III stars could also be ``dark stars'', which would make the detection prospects slightly better (Zackrisson et al. 2010a)\nocite{2010ApJ...717..257Z}. There are also theories where dark stars accrete matter reaching masses of up to 10$^7$~M$_{\odot}$, which would improve prospects for observation significantly (\citealt{2010ApJ...716.1397F}; Zackrisson et al. 2010b\nocite{2010MNRAS.407L..74Z}; \citealt{2012MNRAS.422.2164I}). | 12 | 6 | 1206.0007 |
1206 | 1206.5552_arXiv.txt | We present Keck/LRIS spectra of over 200 galaxies with well-determined redshifts between 0.4 and 1.4. We combine new measurements of near-ultraviolet, low-ionization absorption lines with previously measured masses, luminosities, colors, and star formation rates to describe the demographics and properties of galactic flows. Among star-forming galaxies with blue colors, we find a net blueshift of the \feII\ absorption greater than 200\kms (100\kms) towards 2.5\% (20\%) of the galaxies. The fraction of blueshifted spectra does not vary significantly with stellar mass, color, or luminosity but does decline at specific star formation rates less than roughly $0.8$~Gyr$^{-1}$. The insensitivity of the blueshifted fraction to galaxy properties requires collimated outflows at these redshifts, while the decline in outflow fraction with increasing blueshift might reflect the angular dependence of the outflow velocity. The low detection rate of infalling gas, 3 to 6\% of the spectra, suggests an origin in (enriched) streams favorably aligned with our sightline. We find 4 of these 9 infalling streams have projected velocities commensurate with the kinematics of an extended disk or satellite galaxy. The strength of the \mgII\ absorption increases with stellar mass, B-band luminosity, and U-B color, trends arising from a combination of more interstellar absorption at the systemic velocity {\it and} less emission filling in more massive galaxies. Our results provides a new quantitative understanding of gas flows between galaxies and the circumgalactic medium over a critical period in galaxy evolution. | \label{sec:intro} Key processes in the evolution of galaxies include the flow of cool gas into galaxies, the conversion of these baryons into stars, and the ejection of gas enriched with heavy elements. Determining the factors that govern the circulation of baryons remains a critical, unsolved problem in cosmology. The measured properties of galaxies indicate strong evolution in this baryon cycle between the peak in star-formation activity at redshift $z \approx 2$ and the current epoch (Lilly \et 1996; Madau \et 1996; Hopkins \& Beacom 2006). At any redshift during this period, star-forming galaxies populate a fairly well-defined locus in the star formation rate (SFR) -- stellar mass plane (Bell \et 2005; Elbaz \et 2007; Noeske \et 2007). However, the observed SFR at a fixed stellar mass declines by about a factor of 20 from $z \sim 2$ to the present. The Tully-Fisher relation also evolves in the sense that present-day galaxies have about 2.5 times more stellar mass than do galaxies at $z \approx 2.2$ with the same circular velocity (Cresci \et 2009). Since redshift $z \approx 1$ when roughly half the mass in present-day galaxies was assembled (Drory \et 2005; Faber \et 2007; Marchesini \et 2009), the stellar mass in red sequence galaxies has continued to grow; but the mass in blue galaxies has remained essentially constant (Faber \et 2007). Over this period, some process quenches star formation in massive galaxies, yet observational efforts to identify this process with feedback from active galactic nuclei (AGN) have not reached a consensus (cf., Schawinski \et 2007; Aird \et 2012). The leading explanation for this drop in cosmic star-formation activity is a decline in the accretion rate of cool gas onto galaxies. The simplest form of this idea, that the mean growth rate of dark matter halos regulates the gas inflow rate onto galaxies, fails; it yields too much star formation in low mass halos at early times (Bouch\'{e} \et 2010) and does not shut down star formation in high mass halos at later times (Somerville \et 2008). Introducing gas physics -- heating by virial shocks in massive halos {\it and} outflows in both low-mass and high-mass halos -- effectively tilts the underlying relationship between halo accretion rate and mass (described by the simplest model) into the observed SFR -- stellar mass relation and Tully-Fisher relation (Bouch\'{e} \et 2010 ; Dav\'{e}, Oppenheimer, \& Finlator \et 2011a; van de Voort \et 2011a, b; Dav\'{e}, Finlator, \& Oppenheimer 2012). The balance between gas outflow and inflow then determines the relationship between the gas-phase metallicity and stellar mass (Dav\'{e}, Finlator, \& Oppenheimer 2011b). The problem with this scenario is that this baryon-driven picture of galaxy formation rests on a rather rudimentary understanding of which galaxies have outflows and inflows and how the physical properties of flows vary with fundamental galaxy properties. Gas flows imprint resonance absorption lines on the spectra of their host galaxies that can often be kinematically distinguished from interstellar gas at the systemic velocity (Heckman \et 2000; Schwartz \& Martin 2004; Martin 2005, 2006; Rupke \et 2005; Schwartz \et 2006; Tremonti \et 2007; Martin \& Bouch\'{e} 2009; Rubin \et 2010a). Thus far, studies of gas flows at $0.5 < z < 2$ have either been based on just a few individual galaxies (Sato \et 2009; Coil \et 2011; Rubin \et 2011) or on {\it composite spectra}, i.e., the average of many low S/N ratio spectra (Weiner \et 2009; Rubin \et 2010b). The pioneering study of Weiner \et (2009) concluded that blueshifted \mgII\ absorption was ubiquitous in spectra of star-forming galaxies at $z \sim 1.4$ and demonstrated that the Doppler shift and absorption equivalent width of this absorption increase (rather slowly) with both stellar mass and SFR. At lower redshift, however, Rubin \et (2010) did not find blueshifted \mgII\ absorption in a composite spectrum of galaxies with stellar masses similar to the Weiner \et (2009) sample and speculated that the higher specific SFRs (i.e., SFR per unit stellar mass) characterizing $z \approx\ 1.4$ galaxies might be required to host such outflows. These two studies have demonstrated the existence of outflows and their possible evolution over a key period in the assembly of galaxies but are limited by their reliance on the {\it mean} spectrum of a population. Averaging many spectra together in a composite hides less common features like gas inflows (Sato \et 2009; Rubin \et 2012) and does not allow de-projection of the line-of-sight Doppler shift into an outflow velocity. Here, we present the results of a survey of near-ultraviolet spectral features in 208 galaxies with redshifts between $0.4 < z < 1.4$ to provide an empirical measurement of how outflow and inflow properties change with cosmic time and galaxy properties. The galaxies were selected from the DEEP2 (Deep Extragalactic Evolutionary Probe 2) survey (Davis \et 2003; Newman \et 2012). This redshift survey provides a relatively un-biased sample for investigating the demographics of outflows and inflows in galaxies brighter than $R_{AB} = 24.1$ in four fields. More importantly though, selecting from the DEEP2 survey ensures that fundamental galaxy parameters -- such as stellar mass, $B$-band luminosity, and $U-B$ color -- have been measured for the sample in a systematic manner (Bundy \et 2006; Willmer \et 2006). The additional photometry obtained for the AEGIS (All-Wavelength Extended Groth Strip International Survey) field allows us to measure SFRs for 51 of these galaxies; and these SFRs provide important information about the energy and momentum produced by supernovae, stellar winds, and radiation from massive stars and therefore available to drive outflows (Chevalier \& Clegg 1985; Murray \et 2005). In this paper, we focus on the diagnostics provided by low-ionization, resonance-absorption lines, reserving the presentation of high-ionization interstellar absorption and stellar features, resonance emission, and fluorescent emission for future papers. Low-ionization, resonance lines in near-ultraviolet (rest-frame) spectra provide an empirical bridge between the optical transitions typically observed for low-redshift galaxies and the far-UV transitions studied extensively in spectra of high-redshift galaxies (Steidel \et 2010) because the \mgII\ $\lambda \lambda 2796$, 2803 doublet is accessible from the ground over the broad redshift range from roughly $0.25 < z < 2.5$ (Martin \& Bouch\'{e} 2009; Tremonti \et 2007; Weiner \et 2009; Rubin \et 2010; Coil \et 2011). At the lower end of this redshift range, \mgII\ absorption properties can be cross calibrated with rest-frame optical lines such as \naI\ $\lambda \lambda 5890, 96$ and \ion{Ca}{2} $\lambda 3933, 69$ . At $z \sgreat\ 1.19$, the \mgII\ absorption properties can be directly compared to far-UV transitions in \ion{Si}{2}, \ion{Al}{2}, \ion{C}{2}, and \ion{C}{4} commonly employed to study outflows in much higher redshift galaxies (Shapley \et 2003; Steidel \et 2010; Jones \et 2012). The \mgII\ doublet provides a very sensitive probe of outflows for several reasons: singly-ionized magnesium is a dominant ionization state over a broad range of conditions, both lines have large oscillator strength, and the cosmic abundance of Mg is fairly high. Scattered \mgII\ emission, however, partially fills in the intrinsic absorption troughs in galaxy spectra and complicates the interpretation of the Doppler shift of the absorption trough. Spectral coverage blueward of the \mgII\ $\lambda \lambda 2796$, 2803 doublet provides access to a series of strong \feII\ resonance lines which alleviate the concern about emission filling. The \feII\ absorption troughs provide a cleaner view of the intrinsic absorption profile because fluorescence (rather than resonance emission) often follows absorption in several of these \feII\ transitions (Prochaska \et 2011; Erb \et 2012). Another advantage of \feII\ lines over the \mgII\ doublet arises from the large number of \feII\ transitions in the NUV. The oscillator strengths of the NUV \feII\ transitions span a substantial range and therefore make it possible to place useful bounds on the column density of singly-ionized iron, thereby constraining the total gas columns for an assumed (i.e., model dependent) metallicity and ionization fraction. The paper is organized as follows. Section~\ref{sec:observations} introduces the sample and describes our new Keck observations, including a discussion of the systemic velocity determination and absorption-line sensitivity. After providing a broad overview of the near-UV spectral features, Section~\ref{sec:diagnostics} explains the complications caused by emission filling and how we identify and measure galactic gas flows. We then compare the line profiles of the \mgII\ and \feII\ absorption troughs in Section~\ref{sec:compare}. In Section~\ref{sec:galaxy_properties}, we calculate the fraction of galaxies with a Doppler shift of low-ionization absorption relative to the systemic velocity and use previously measured galaxy properties to illustrate the demographics of the galaxies showing blueshifts. We then discuss the physical properties of the outflows and show how the outflow properties scale with galaxy properties. In Section~\ref{sec:inflow}, we expound on our discovery of net inflows of enriched gas. Section~\ref{sec:conclusions} summarizes our conclusions. Throughout this paper we assume a cosmological model with $\Omega_m = 0.3$, $\Omega_{\Lambda} = 0.7$, and $H_0 = 70$\kms\ Mpc$^{-1}$. We adopt the atomic oscillator strengths and cosmic abundance ratios given by Morton (2003) as well as the associated vacuum wavelengths for transitions shortward of 3200\AA. We refer to optical transitions by their wavelengths in air (for ease of comparison to previously published work), but we work with their vacuum wavelengths. \clearpage | Many properties of the infalling gas are consistent with the cold flow scenario. The velocity components along our sightlines are largely consistent with virial motion, the exception being the extremely large redshift towards 22005270 which is associated with the satellite galaxy 22005066. The two-dimensional spectra provide spatial information about the inflow in 3 other objects, and in those the redshift of the infalling gas matches the projected velocity on one side of a rotating, gas disk. This association could reflect a stream crossing in front of each galaxy and connecting to the outer disk in the manner predicted by Stewart \et (2011a, b). The low incidence of redshifted \feII\ absorption in spectra of $z \sim 1$ galaxies would also be expectd for cold flows. To have any chance of detecting a cold flow in a metal line, the sightline must run along the filament (e.g., see Figure~2 in Kimm \et 2010), and the chance of such a favorable orientation is small. While differences in the radiative transfer calculations produce some variation in the average covering fraction among these models, galaxies with covering fractions of more than a few percent, as originally suggested by Dekel \et 2009a, are now thought to be quite rare (Kimm \et 2009). Numerical simulations suggest that only a few percent of the sightlines passing within 100 kpc of a massive galaxy at $z \sim 2$ intersect dense ($\log N_H (\col) > 20.3$), infalling gas (Kimm \et 2010; Faucher-Gigu\`{e}re \et 2011; Fumagalli \et 2011). It must be acknowledged, however, that these high columns represent the cool gas threading the filaments rather than the smooth component of the cold streams, which dominates the cross section of neutral hydrogen absorption below $\log N_{HI} (\col) = 18$ within $R_{vir}$ (Fumagalli \et 2011). At $z < 2$, less than half the cross section in the range $19 < \log N_{HI} (\col) < 20$, and less at higher columns, is likely due to streams (Fumagalli \et 2011). Theory predicts the disappearance of cold flows in halos more massive than roughly $10^{12}$\msun. The halo masses of many of the galaxies in the high stellar mass tertile of the LRIS sample are likely above $10^{12}$\msun\ based on either their clustering properties (Coil \et 2008) or halo abundance matching (Behroozi, Conroy, and Wechsler 2010). In the LRIS sample, we did not find any change in the inflow fraction with increasing stellar mass; however, this sample excludes red sequence galaxies. As can be seen by comparing the density of red sequence and blue cloud galaxies in \fig~\ref{fig:sample}, the fraction of red galaxies increases quite quickly among galaxies with stellar masses $\log M/ \msun > 10.3$, the typical stellar mass in a $10^{12}$\msun\ halo at $z \sim 1$ indicated by abundance matching (Behroozi \et 2010). Whether or not the fraction of inflow galaxies declines with mass in an unbiased sample cannot be determined from our data since we do not measure the inflow fraction towards red galaxies. Likewise, the lack of any evolution with redshift might be limited to blue cloud galaxies which compose the LRIS sample. According to some models (Kimm \et 2011; Fumagalli \et 2011), detection of the inflows via a metal line definitely implies these are not cold accretion flows. From our perpsective, however, the metallicity of the cold inflows near galaxies remains a highly uncertain quantity in the models due to the unrealistic treatment of metal recycling via galactic winds. The ciculation of metals might significantly enrich the cold flows before the infalling gas is incorporated into the galactic disk. In this recycling scenario, the distinction between cold flows and infalling relic outflows may become blurred. Numerical simulations including outflows have suggested that the covering factor of high column density (DLA-like) gas increases with halo mass, due in large part to the kinematics of outflowing gas (Hong \et 2010; Faucher-Gigu\`{e}re \et 2011). Given the lack of a strong scaling relation between outflow velocity and stellar mass, we expect a larger fraction of the outflowing material to fall back onto the more massive galaxies. Hence, the recycling of outflows should produce an increasing inflow column as stellar mass increases. The distinction between the increasing (decreasing) inflow column with stellar mass due to wind recycling (cold flows) may provide a means to discriminate the primary origin of infalling metal-enriched gas. This proposal further emphasizes that inflow fraction measurements should include red sequence galaxies. Although we expect only star-forming galaxies to drive outflows, the galaxies may migrate across the color - magnitude diagram (due to the cessation of star formation) before the bulk of the outflow slows down and turns around (due to the gravitational attraction of the galaxy). The first inflow galaxies identified appeared to support this scenario; their optical colors placed them on the red sequence, but their UV-optical colors indicated they were forming stars a few 100 Myr prior (Sato \et 2009). A subsequent study of 13 K+A galaxies, however, discovered 2 inflow galaxies (Coil \et 2011) suggesting an inflow fraction similar to our result for blue cloud galaxies.\footnote{ The galaxies in our LRIS sample that show redshifted \feII\ absorption do not have the spectral signatures of K+A galaxies nor the the colors of green-valley galaxies. They are normal blue cloud galaxies and are not post-starburst galaxies.} Finally, mergers also bring gas into galaxies. The association of the galay 22005066 with the redshifted \feII\ absorption in 22005270 may reflect this scenario. Because 22005066 is the less luminous of the two galaxies, we would expect a tidal stream to pull gas out of 22005066 onto 22005270. In summary, cold flows, mergers, and wind recycling are all expected to contribute to gas infall. The low incidence of infall suggests that whichever process dominates, the covering fraction of the infalling streams is low. Much larger samples of individual spectra will be needed to further distinguish among these mechanisms on a statistical basis. | 12 | 6 | 1206.5552 |
1206 | 1206.2144_arXiv.txt | Using high time cadence images from the STEREO EUVI, COR1 and COR2 instruments, we derived detailed kinematics of the main acceleration stage for a sample of 95 CMEs in comparison with associated flares and filament eruptions. We found that CMEs associated with flares reveal on average significantly higher peak accelerations and lower acceleration phase durations, initiation heights and heights, at which they reach their peak velocities and peak accelerations. This means that CMEs that are associated with flares are characterized by higher and more impulsive accelerations and originate from lower in the corona where the magnetic field is stronger. For CMEs that are associated with filament eruptions we found only for the CME peak acceleration significantly lower values than for events which were not associated with filament eruptions. The flare rise time was found to be positively correlated with the CME acceleration duration, and negatively correlated with the CME peak acceleration. For the majority of the events the CME acceleration starts before the flare onset (for 75\% of the events) and the CME accleration ends after the SXR peak time (for 77\% of the events). In $\sim$60\% of the events, the time difference between the peak time of the flare SXR flux derivative and the peak time of the CME acceleration is smaller than $\pm$5 min, which hints at a feedback relationship between the CME acceleration and the energy release in the associated flare due to magnetic reconnection. | Solar flares and coronal mass ejections (CMEs) are the two most energetic phenomena in our solar system. Solar flares are abrupt releases of energy up to 10$^{25}$~J within tens of minutes and can be observed in the whole electromagnetic spectrum from radio emission to $\gamma$-rays \citep[e.g.][]{fletcher2011}. CMEs are sporadic ejections of coronal material with velocities in the range of $\sim$100--3000~km~s$^{-1}$ \citep[e.g.][]{yashiro2004, gopalswamy2009}. It is generally accepted that both CMEs and flares are different manifestations caused by magnetic reconnection in the solar corona, but the details how both phenomena are related are still under investigation. Various statistical studies using white light coronagraphic observations showed positive correlations between the flare intensity and CME velocity \citep[e.g.][]{moon2002, vrsnak2005, yashiro2009} or CME kinetic energy \citep{burkepile2004, yashiro2009}. The temporal differences between the CME and the flare onset were found to be quite small. \citet{michalek2009} and \citet{yashiro2009} found a Gaussian distribution for the difference between the flare and CME onsets. According to \citet{michalek2009} the mean difference is 7 min. In these studies, linear back extrapolation of the CME height time curve was used to estimate the CME start time. However, the CME onset times cannot be accurately determined by using only coronagraphic observations which miss the early phase of the CME acceleration and propagation due to the occulter disk. The CME kinematics typically shows three phases of evolution \citep{zhang2001}. In the initiation phase, the CME rises with velocities of several tens of km~s$^{-1}$, followed by an impulsive acceleration. Thereafter, the CME propagates with almost constant or slowly decreasing/increasing velocity depending on its interaction with the ambient solar wind \citep[e.g.][]{gopalswamy2000}. Recent case studies \citep[e.g.][]{gallagher2003, vrsnak2004, temmer2008} combined EUV images with coronagraphic observations to derive detailed CME acceleration profiles from the CME initiation close to the solar surface until its propagation beyond 15~$R_{\odot}$. In these events the CME acceleration profiles showed a good synchronization with the energy release in the associated flare, as evidenced in the HXR flux or SXR derivative. \citet{maricic2004} studied a set of 22 CME events associated with flares using SOHO EIT, MLSO Mark IV, LASCO C2 and C3. They report correlations between various CME and flare parameters as well as a close synchronization between the CME acceleration and the flare SXR flux derivative for $\geq$ 50\% of the events under study. In \citet[][to which we refer to as paper I in the following]{bein2011}, we presented statistics and correlations between various decisive CME parameters for a sample of 95 events: peak velocity, peak acceleration, acceleration duration, height at peak velocity, height at peak acceleration and initiation height. To this aim we combined EUV images from the Extreme Ultraviolet Imager \citep[EUVI;][]{wuelser2004} with COR1 and COR2 coronagraphic observations onboard the STEREO mission. The high time cadence and the overlapping field of views (FOV) of the different STEREO instruments enabled us to derive detailed and continuous CME height-, velocity- and acceleration-time profiles from the low corona up to about 15~R$_{\odot}$. Out of the 95 CMEs presented in \citet{bein2011}, 70 events could be associated with a GOES (Geostationary Operational Environmental Satellite) flare and 24 events with a filament eruption. 9 events are associated with both, a flare and a filament. In the present paper, we perform a statistical study on the relation between characteristic CME and flare parameters as well as on the temporal relationship between the two phenomena. We also study the characteristic CME parameters separately for events with/without associated flares and with/without associated filament eruptions. | Based on a sample of 95 CME events we present a statistical study on various characteristic CME parameters and their relation to flares. CMEs which are associated with flares show on average higher peak velocities ($v_{max}$), higher peak accelerations ($a_{max}$), shorter acceleration phase durations ($t_{acc}$), lower heights at peak velocity ($h_{vmax}$), lower heights at peak acceleration ($h_{amax}$) and lower initiation heights ($h_{0}$). The ratio between the median values of the lognormal probability functions of both subgroups is about a factor of 2. Only for $v_{max}$ the ratio is significantly smaller ($\sim$1.1), most probably due to the small range of $v_{max}$ values. Due to the anticorrelation between $a_{max}$ and $t_{acc}$ and the relation $v_{max}\approx a_{max}\cdot t_{acc}$ the range for $v_{max}$ is smaller than for the other two quantities, it basically covers only one order of magnitude. Although we found clear differences in the mean and median values of the two subgroups there exist also events associated with flares, which have low $v_{max}$ and $a_{max}$ values and high $t_{acc}$, $h_{vmax}$, $h_{amax}$ and $h_{0}$. For instance, the smallest measured $a_{max}$ value for a CME with flare was 77~m~s$^{-2}$ and the highest measured value for a CME without flare was 1577~m~s$^{-2}$. The Kolmogorov-Smirnov test suggested a distinction between both distributions (flare/non-flare associated events) for every CME parameter (except for $v_{max}$) at 0.05--0.15 levels of significance. The clearest distinctions were found for $a_{max}$ and $h_{amax}$ at a 0.05 level of significance. CMEs which are associated with erupting filaments show on average smaller $v_{max}$, $a_{max}$, and larger $t_{acc}$, $h_{vmax}$, $h_{amax}$ and $h_{0}$. These trends are the other way round than for the flare association. The ratio between the mean and median values of both subgroups is somewhat smaller than for the flare association. Again we found the smallest ratio between the median values $\mu^*$ of $v_{max}$ (1.1). The $\mu^*$ of $a_{max}$ and $h_{vmax}$ showed the highest ratio with 1.9. For $t_{acc}$, $h_{amax}$ and $h_{0}$ we found ratios between 1.2 and 1.5. But there exist also CMEs associated with filament eruptions, which have high $v_{max}$ and $a_{max}$ values and low $t_{acc}$, $h_{vmax}$, $h_{amax}$ and $h_{0}$. For example the highest measured $a_{max}$ value for a CME which was associated with an erupting filament was 1561~m~s$^{-2}$, the smallest value for a CME event with no erupting filament was 35~m~s$^{-2}$ (the second lowest value of the whole distribution). Both of these events were not associated with a flare. The Kolmogorov-Smirnov test did not show a clear distinction between both subgroups. Only for $a_{max}$ the test suggested that both distributions do not come from the same population at a 0.05 level of significance. The correlations obtained between $v_{max}$, $a_{max}$ and $t_{acc}$ with the GOES peak flux $F_{SXR}$ and the SXR rise time $t_{SXR}$ of the associated flare were low. We found a weak positive correlation between the CME acceleration duration $t_{acc}$ and the flare rise time $t_{SXR}$ ($c=0.37\pm0.15$) and a weak negative correlation between $a_{max}$ and $t_{SXR}$ ($c=-0.32\pm0.15$). Correlations between $v_{max}$ and $a_{max}$ with $F_{SXR}$ were $c=0.32\pm0.13$ and $c=0.28\pm0.12$ respectively. If the events are averaged and binned into the different GOES classes, the correlations are much more distinct. The mean values in each GOES class show an increasing trend for $v_{max}$ and $a_{max}$, and a decreasing trend for $t_{acc}$ with higher GOES flux. For the majority of the events (75\%) we found that the CME acceleration starts \textit{before} the SXR flare onset, which is consistent with the findings of \citet{maricic2007}, suggesting that the flare is a consequence of the eruption. Similar to our study these authors also found about one fourth of the events, for which the onset of the associated flare occurred before the CME acceleration started. They explained these cases by a superposition of two flares, a confined flare in the pre-eruption stage, which releases only a part of the stored magnetic energy, and a second flare, beginning after the CME acceleration onset and causing a prolongation of the first flare in the full-disk integrated SXR light curve. The CME is associated with the second flare but because of the superposition, the flare start is measured from the first one. To test this hypothesis, we checked the SXR curves for all events, for which the CME acceleration start was after the flare onset and found indeed in 11 out of 14 events evidence of a second SXR peak, confirming their suggestion \footnote{This is a significantly higher rate than in the total sample of events, in which about 20\% showed a double SXR peak}. Figure \ref{accsxr} shows two representative examples. The top panels show the CME acceleration profile, the bottom panels the GOES SXR flux together with its derivative. We marked the two possible SXR peaks by arrows. Assuming that we have a superposition of two subsequent flares, we would also measure erroneous SXR start times and as a result too long $t_{SXR}$. This misidentification would also influence our correlation negatively. To test this, we correlated $t_{SXR}$ again with $a_{max}$ and $t_{acc}$ considering only events for which the CME acceleration starts before the flare onset and found indeed distinctly higher correlation coefficients (Figure \ref{risetime1}, $c=0.59\pm0.12$ and $c=-0.50\pm0.14$) than for the whole sample (shown in Figure \ref{risetime}, $c$=0.37$\pm$0.15 and $c=-0.32\pm0.15$). Thus superposition of two subsequent flares and more complex structures are probably a reason for the weak correlations and may account for a considerable number of events, where the flare seems to start before the eruption. For the majority of the events (77\%) we found that the end of the CME acceleration occurred \textit{after} the SXR peak. Especially long duration flares reach a certain point, when they become too weak to compensate cooling of the hot plasma. As a result the SXR curve decreases although the energy release in the flare may still be going on. For 81\% of the events the time delay between the CME acceleration peak and the peak of the GOES SXR flux derivative, which is a proxy for the flare energy release rate, was smaller than $\pm$10 min, for 58\% smaller than $\pm$5 min. This high synchronization hints at a feed-back relationship between the CME and the flare energy release \citep{lin2004, zhang2006, maricic2007, temmer2008, reeves2010, temmer2010}. There are basically two forces acting on a flux rope in equilibrium, an upward directed magnetic pressure and a downward directed magnetic tension of the overlying magnetic field. When the magnetic structure looses equilibrium, it starts rising and a current sheet is formed below it, where magnetic reconnection takes place \citep{priest2002}. The reconnection reduces the tension of the overlying field and enhances the magnetic pressure at the bottom part of the flux rope due to additional poloidal flux, providing the upward acceleration of the rope \citep{vrsnak2008}. The upward motion of the rope leads to elongation of the current sheet and a more efficient reconnection, thus enhancing the acceleration. On the other hand, more efficient reconnection means also a more powerful energy release in the CME-associated flare, which directly relates the dynamics of the eruption and the energy release in the flare. \begin{figure*} \centering \includegraphics[scale=0.7]{fig1.eps}~~~~~~~~~~~~ \includegraphics[scale=0.7]{fig2.eps} \caption{CME kinematics and GOES 1--8 \AA~soft X-ray flux for the CME-flare events that occurred on 2008 January 7 (left) and 2009 January 9 (right). The top panels show the measured CME height-time curve derived from STEREO EUVI (crosses), COR1 (triangles) and COR2 (asterisks) observations. The measurement errors (0.03 $R_{\odot}$ for EUVI, 0.125 $R_{\odot}$ for COR1 and 0.3 $R_{\odot}$ for COR2), which are in some cases smaller than the plot symbols, and the spline fit (solid line) are overplotted. The second and third panels show the CME velocity and acceleration profiles, derived from numerical differentiation of the direct measurements (plot symbols) and the spline fit (solid line) to the height-time curve. The grey shaded area represents the error range of the spline fit. The bottom panels show the GOES flux (black solid line) and its derivative (red dashed line) of the associated flare.} \label{height1} \end{figure*} \begin{figure*} \centering \includegraphics[scale=0.7]{fig3.eps}~~~~~~~~~~~~ \includegraphics[scale=0.7]{fig4.eps} \caption{Same as Figure \ref{height1} but for the events observed on 2009 December 22 (left) and 2010 February 12 (right).} \label{height2} \end{figure*} \begin{figure} \centering \includegraphics[scale=1.1]{fig5.eps} \caption{Histogram of the CME peak velocity $v_{max}$ for the whole sample of 95 events (grey distributions in the top and bottom panels). In the top panel the $v_{max}$ distribution of CME events associated with a flare is overlaid in color, in the bottom panel the histogram of CME events associated with an erupting filament is overlaid.} \label{histvmax} \end{figure} \begin{figure} \centering \includegraphics[scale=1.1]{fig6.eps} \caption{Same as Fig. \ref{histvmax} but for the CME peak acceleration.} \label{histamax} \end{figure} \begin{figure} \centering \includegraphics[scale=1.1]{fig7.eps} \caption{Same as Fig. \ref{histvmax} but for the CME acceleration duration.} \label{histaccdur} \end{figure} \begin{figure} \centering \includegraphics[scale=1.1]{fig8.eps} \caption{Same as Fig. \ref{histvmax} but for the height $h_{0}$ where the CME could be first identified. This height is used as an estimate of the CME initiation height.} \label{histh0} \end{figure} \begin{figure} \centering \includegraphics[scale=1.1]{fig9.eps} \caption{Same as Fig. \ref{histvmax} but for the height $h_{vmax}$, where the CMEs reached their maximum velocity. } \label{histhv} \end{figure} \begin{figure} \centering \includegraphics[scale=1.1]{fig10.eps} \caption{Same as Fig. \ref{histvmax} but for the height $h_{amax}$, where the CMEs reached their maximum acceleration. } \label{histha} \end{figure} \begin{figure} \centering \includegraphics[scale=0.5]{fig11.eps} \caption{Distribution of the GOES classes of the flare events under study.} \label{goesclass} \end{figure} \begin{figure} \centering \includegraphics[scale=1]{fig12.eps} \caption{CME peak velocity (top) and peak acceleration (bottom) against the GOES peak flux of the associated flare. A double logarithmic space is used for the plot and the calculation of the correlation coefficient $c$, which is annotated in each panel. The regression line is overplotted in red.} \label{goesmax} \end{figure} \begin{figure} \centering \includegraphics[scale=1]{fig13.eps} \caption{CME acceleration duration (top) and peak acceleration (bottom) against the flare SXR rise time. The correlation coefficient $c$ (calculated in double logarithmic space) and the regression line is overplotted. } \label{risetime} \end{figure} \begin{figure} \centering \includegraphics[scale=0.7]{fig14.eps} \caption{Distribution of the risetime $t_{SXR}$ of the GOES SXR flares.} \label{histrisetime} \end{figure} \begin{figure} \centering \includegraphics[scale=1]{fig15.eps} \caption{Mean value of the CME peak velocity $v_{max}$ (top), acceleration $a_{max}$ (middle) and acceleration duration $t_{acc}$ (bottom) for each GOES class plotted against the mean GOES flux in each subgroup. The correlation coefficient $c$ (calculated in logarithmic space) is annotated in each panel.} \label{goessep} \end{figure} \begin{figure*} \centering \includegraphics[scale=0.75]{fig16.eps} \caption{Distribution of the time differences between the start of the GOES flare and the start of the CME acceleration (top), GOES peak time and CME acceleration end time (middle) and the peak time of the derivative of the GOES 1--8 \AA~flux and the CME acceleration peak time (bottom). On the left hand side, the time differences are plotted in minutes, whereas on the right hand side the time differences are normalized by the acceleration duration of the corresponding CMEs. Positive values mean that the CME acceleration start, end or peak occurred before the flare SXR start, peak or derivative peak, respectively.} \label{histtime} \end{figure*} \begin{figure*} \centering \includegraphics[scale=0.75]{fig17.eps} \caption{CME acceleration (top) and GOES SXR curve (bottom) from the events observed on 25 October 2009 (left hand side) and 23 March 2008 (right hand side). In both events the CME acceleration starts after the flare onset, which may be related to an overlap of two flares. Arrows mark two SXR maxima indicating two possible flare maxima.} \label{accsxr} \end{figure*} \begin{figure} \centering \includegraphics[scale=1.0]{fig18.eps} \caption{Same as Figure \ref{risetime} but only for events, where the CME acceleration starts before the flare.} \label{risetime1} \end{figure} | 12 | 6 | 1206.2144 |
1206 | 1206.3886_arXiv.txt | We present the results of a photometric and astrometric study of the low mass stellar and substellar population of the young open cluster Blanco 1. We have exploited $J$ band data, obtained recently with the Wide Field Camera (WFCAM) on the United Kingdom InfraRed Telescope (UKIRT), and 10 year old $I$ and $z$ band optical imaging from CFH12k and Canada France Hawaii Telescope (CFHT), to identify 44 candidate low mass stellar and substellar members, in an area of 2 sq. degrees, on the basis of their colours and proper motions. This sample includes five sources which are newly discovered. We also confirm the lowest mass candidate member of Blanco 1 unearthed so far (29M$_{\rm Jup}$). We determine the cluster mass function to have a slope of $\alpha$=+0.93, assuming it to have a power law form. This is high, but nearly consistent with previous studies of the cluster (to within the errors), and also that of its much better studied northern hemisphere analogue, the Pleiades. | Open clusters are often acclaimed as excellent laboratories with which to study star formation. This is due to the co-eval nature of their members and estimates of their age being comparatively robust. Many open star clusters have been studied to date, yielding a large number of low mass members (e.g. \citealt{Baker_2010,Casewell_2007,Lodieu_2007a}) which have been used to refine our knowledge about the low mass end of star formation via mapping the initial mass function (IMF). The IMF, the number of objects per unit mass interval, is an observable outcome of star formation and can be used to critically examine theoretical models of this process.The IMF is commonly measured using an $\alpha$ parameter, where dN/dM $\propto$ M$^{-\alpha}$ and N is number of objects, and M is mass. For most open star clusters (ages ~100 Myr), $\alpha$ is roughly consistent across all samples and $\approx$0.6 \citep{bouvier05}. This value is also consistent with. field values such as those of \citet{Chabrier_2003}, although recently it has been suggested that for very low mass field brown dwarfs the IMF may have a different form. Indeed \citet{burningham10} suggest that in this case $\alpha$ may even have a negative value. In recent years there has been a particular emphasis on building a solid comprehension of the mechanisms by which very low mass brown dwarfs and free-floating planetary mass objects form (e.g. \citealt{bate11}). Nevertheless, key questions remain to be answered e.g. what is the lowest possible mass of object that can be manufactured by the star formation process$?$ From a theoretical stance, traditional models predict that if substellar objects form like stars, via the fragmentation and collapse of molecular clouds, then there is a strict lower mass limit to their manufacture of 0.007-0.010 $\Msun$. This is set by the rate at which the gas can radiate away the heat released by the compression (e.g. \citealt*{low76}). However, in more elaborate theories, hypothetical magnetically mediated rebounds in collapsing cloud cores might lead to the decompressional cooling of the primordial gas, a lowering of the Jeans mass and hence the production of gravitationally bound fragments with masses of only $\sim$0.001 $\Msun$ \citep{Boss_2001}. However, while many surveys of open star clusters have been performed to search for substellar members, the majority of these are in the heavily populated Northern hemisphere clusters. The lack of southern coverage from surveys (e.g. Sloan Digital Sky Survey \citealt{York_2000}; UKIRT Infrared Deep Sky Survey \citealt{Warren_2007}) has impeded detailed studies of the substellar population of a plethora of potentially interesting southern open clusters. Blanco 1 is a 90$\pm$25 Myr \citep{Panagi_1997} open cluster with an age similar to that of the 125 Myr Pleaides cluster \citep*{Stauffer_1998} at a distance of 207$\pm$12 pc as determined from $Hipparcos$ measurements \citep{Leeuwen_2009}. Recent work on the cluster includes spectroscopy of F and G type stars \citep{Ford_2005} which show that the metallicity is [Fe/H]=+0.04, with subsolar abundances for [Ni/Fe], [Si/Fe], [Mg/Fe], and [Ca/Fe]. \citet*{cargile10} have determined a Lithium age for the cluster of 132$\pm$24 Myr which is closer to the age of the Pleiades than that of \citet{Panagi_1997}. We have taken the age of the cluster to be 120 Myr which is close to both measured values, and is present in the \citet{Chabrier_2000} DUSTY models. Recently \citet{platais11} surveyed 11 square degrees of the cluster to provide a comprehensive proper motion catalogue for all stellar objects down to M5V. \citet{Moraux_2007} performed the first study of the cluster to search for brown dwarfs using CFH12k on the Canada-France-Hawaii Telescope in the optical $z$ and $I$ bands to image 2.3 square degrees of the cluster centre. They discovered $\approx$ 300 cluster members; 30-40 were estimated to be brown dwarfs, some of which had additional $K$ band photometry and optical spectroscopy. Three of these objects were subsequently confirmed as members by \citet{cargile10}. We have used the $I$ and $z$ band images from \citet{Moraux_2007} and have combined them with additional deep ($J\approx22$) $J$ band photometry obtained using WFCAM on UKIRT allowing us to not only select fainter candidate cluster members, but also to measure the proper motion for some of the previously identified objects to prove if they are indeed associated with the cluster. | We have used near-IR and optical photometry with proper motions derived from CFHT $z$ and WFCAM $J$ band images to identify 44 candidate cluster members with masses between 29 M$_{\rm Jup}$ and 80 M$_{\rm Jup}$. 5 of these are previously unidentified candidate members and 40 have been identified by \citealt{Moraux_2007}, 8 of which have been confirmed as cluster brown dwarfs from spectra. We derive $\alpha$=0.93$\pm$0.11 from the mass spectrum, which is consistent with the literature for this cluster. | 12 | 6 | 1206.3886 |
1206 | 1206.6147_arXiv.txt | We have observed 3C~279 in a $\gamma$-ray flaring state in November 2008. We construct quasi-simultaneous spectral energy distributions (SEDs) of the source for the flaring period of 2008 and during a quiescent period in May 2010. Data have been compiled from observations with Fermi, Swift, RXTE, the VLBA, and various ground-based optical and radio telescopes. The objective is to comprehend the correspondence between the flux and polarization variations observed during these two time periods by carrying out a detailed spectral analyses of 3C~279 in the internal shock scenario, and gain insights into the role of intrinsic parameters and interplay of synchrotron and inverse Compton radiation processes responsible for the two states. As a first step, we have used a multi-slice time-dependent leptonic jet model, in the framework of the internal shock scenario, with radiation feedback to simulate the SED of 3C~279 observed in an optical high state in early 2006. We have used physical jet parameters obtained from the VLBA monitoring to guide our modeling efforts. We briefly discuss the effects of intrinsic parameters and various radiation processes in producing the resultant SED. | Introduction} Blazars are well known for their variability and power of polarized radiation across a wide range of the electromagnetic spectrum \citep[]{df2007, js2007, js2005}. In some cases, the flux can vary on timescales as short as an hour or less \citep[see e.g.,][]{ga1996, at2007}. Blazars exhibit a doubly-peaked spectral energy distribution (SED), in which the low-energy component could extend from radio through UV or X-rays while the high-energy component extends from X-rays to $\gamma$-rays. The low-frequency component of the SED is almost certainly due to synchrotron emission from nonthermal, ultra-relativistic electrons. The high-frequency component, on the other hand, is a result of inverse Compton scattering of seed photons by the same ultra-relativistic electrons producing synchrotron emission (in a leptonic jet scenario). In this case, the seed photons could be the synchrotron photons produced within the jet (synchrotron self Compton, SSC) \citep[]{mg1985, gm1998}, and/or external photons entering the jet from outside (the EC process) \citep[e.g.,][and references therein]{bm2007}. The spectral variability patterns and SEDs are key ingredients in determining the acceleration of particles and the time-dependent interplay of various radiation mechanisms responsible for the observed emission. Another defining characteristic of blazars is the high degree of linear polarization at optical wavelengths. Many bright $\gamma$-ray blazars that are in the \textit{Fermi-LAT} Bright $\gamma$-Ray Source List \citep{aa2009} have shown spectral and linear polarization variability \citep[]{ma2010, df2009, gd2006}. Linear polarization at mm, IR, and optical wavelengths tends to exhibit similar position angles and sometimes correlation across these wavebands, often with some time delay \citep[]{df2009, js2007, ls2000, gss1996, gs1994}. Such correspondence between the variation in polarization and flux across a wide range of the electromagnetic spectrum, combined with VLBI imaging, can be used to identify the location of variable emission at all wavebands and shed light on the physical processes responsible for the variability \citep{ma2010}. The blazar 3C~279, located at a redshift of 0.538 \citep{br1965}, is one of the most prominent and well-studied blazars. This is due to its highly variable nature (change in magnitude $\Delta m \sim 5$ at optical bands) at all wavelengths and high optical polarization up to 45.5\% observed in the U band \citep{me1990}. Intensive multiwavelength campaigns \citep[see, e.g.,][]{lv2008, cr2008, co2007} and theoretical efforts \citep[e.g.,][]{bp2009, bm1996} have led to some important conclusions about the physical properties of 3C~279. Chatterjee et al. (2008) showed that the flux variability in 3C~279 has been found to be significantly correlated at X-rays, optical R band, and 14.5 GHz wave bands, which also suggests that nearly all X-rays are produced in the jet. The X-ray flux has also been associated with superluminal knots, as suggested by correlation with the flux of the core region in the 43 GHz VLBA images \citep{cr2008}. Nevertheless, the nature and origin of its high-energy emission and the relationship of its behavior to the physical aspects of the jet remain elusive \citep[]{lv2008, co2010}. In addition, the correspondence between high- and low-energy emission is also not very well understood. Here, we aim to understand the physical state of 3C~279 at different flux levels and look for a correspondence between the flux and polarization variation observed during the flaring state of November 2008 and quiescent state of May 2010. This can be achieved by carrying out a detailed spectral analyses of 3C~279 under the internal shock scenario, and understanding the role of intrinsic parameters and interplay of synchrotron and inverse Compton radiation processes in shaping the corresponding spectra of the two states. As a first step toward comparing the physical state of 3C~279 at different flux levels, we use the multi-slice time-dependent leptonic jet model of Joshi \& B\"ottcher (2011) (in the framework of internal shock scenario) with radiation feedback to simulate the SED of 3C~279 corresponding to the optical high state of early 2006. The broadband emission of 3C~279 for this state indicates suppressed external Compton emission, which makes it an ideal candidate for simulation using a synchrotron-SSC model. We use physical jet parameters obtained from the VLBA monitoring to guide our modeling efforts and discuss the role of various intrinsic parameters and radiation processes in producing the resultant SED. We briefly describe the model of Joshi \& B\"ottcher (2011) in \S \ref{model}. We discuss our findings about the connection between the flux and polarization variation observed in 3C~279 during flaring and quiescent periods of November 2008 and May 2010, respectively, in \S \ref{3c2}. We discuss our first results from this study in \S \ref{results}. We summarize and give a brief description of future work in \S \ref{summary}. | 12 | 6 | 1206.6147 |
|
1206 | 1206.4697_arXiv.txt | \noindent We study the properties of low--velocity material in the line of sight towards nearby Type Ia Supernovae (SNe Ia) that have measured late phase nebular velocity shifts ($v_{\rm neb}$), thought to be an environment--independent observable. We have found that the distribution of equivalent widths of narrow blended Na I D1 \& D2 and Ca II H \& K absorption lines differs significantly between those SNe Ia with negative and positive $v_{\rm neb}$, with generally stronger absorption for SNe Ia with $v_{\rm neb} \ge 0$. A similar result had been found previously for the distribution of colors of SNe Ia, which was interpreted as a dependence of the temperature of the ejecta with viewing angle. Our work suggests that: 1) a significant part of these differences in color should be attributed to extinction, 2) this extinction is caused by an asymmetric distribution of circumstellar material (CSM) and 3) the CSM absorption is generally stronger on the side of the ejecta opposite to where the ignition occurs. Since it is difficult to explain 3) via any known physical processes that occur \emph{before} explosion, we argue that the asymmetry of the CSM is originated \emph{after} explosion by a stronger ionizing flux on the side of the ejecta where ignition occurs, probably due to a stronger shock breakout and/or more exposed radioactive material on one side of the ejecta. This result has important implications for both progenitor and explosion models. | Type Ia Supernovae (SNe Ia) are important tools for understanding the evolution of the Universe because of their high luminosities and light--curve homogeneity, which led to their standardization for cosmological distance measurements \citep{1993ApJ...413L.105P, 1996AJ....112.2408H} and the discovery of the acceleration of the Universe \citep{riess98, perl99}. They are also important because of their complex nucleosynthetic output and high ejecta kinetic energies ($10^{51}$ erg), which make them key ingredients for galaxy evolution theory. Unfortunately, we still lack a clear understanding of the nature of SN Ia progenitors. Although it is generally accepted that their progenitors are carbon oxygen white dwarfs (CO WD) in mass transfering binary systems \citep[for a review, see][]{2000ARA&A..38..191H}, there is no agreement on what their companions are, whether the CO WD reaches the Chandrasekhar mass at ignition, what the mechanism that triggers the ignition is and how the CO WD burns to form the ejecta. Perhaps the two scenarios that best match observed SN Ia spectra and light curves are those considered in \citet{2012ApJ...750L..19R}: either a stably accreting carbon oxygen white dwarf (CO WD) that reaches the Chandrasekhar mass ($M_{\rm Ch}$), undergoes a thermonuclear runaway and burns in a deflagration to detonation transition, in the $M_{\rm Ch}$--single degenerate (SD) scenario \citep{1984ApJ...286..644N, 1991A&A...245L..25K}; or a pair of CO WDs that merge and ignite dynamically in the violent double degenerate (DD) scenario \citep{2012ApJ...747L..10P}. Two significant differences between both scenarios are the typical central densities at ignition and the distribution of circumstellar material (CSM) around their progenitor systems. In the $M_{\rm Ch}$--SD scenario the central density at ignition is sufficiently high to avoid the synthesis of radioactive $^{56}$Ni in favor of stable iron group elements (IGEs) in the innermost regions of the ejecta. The $M_{\rm Ch}$--SD scenario has potentially abundant CSM that could be produced by either the wind of the donor star, weak nova explosions experienced by the accretor, or by mass loss due to inefficient mass transfer. In the violent DD--scenario, conversely, the primary central density is never high enough to avoid the synthesis of radioactive $^{56}$Ni in the innermost regions of the ejecta, and since the merger and ignition happen in a dynamical time--scale there is no time to leave a significant imprint in the CSM. The $M_{\rm Ch}$--SD scenario appears to over--produce IGEs \citep{1984ApJ...286..644N}, which are not produced significantly in the violent DD scenario. Observations of flat--topped profiles of Fe II lines in the optical and NIR at late phases has been argued to be evidence for stable IGEs in the innermost regions of the ejecta \citep{2004ApJ...617.1258H, 2007ApJ...661..995G}. Moreover, significant nebular velocity shifts ($v_{\rm neb}$) measured in these flat--topped profiles have been suggested as evidence for off--center ignition scenarios, which is expected in $M_{\rm Ch}$--SD scenarios \citep[see e.g.][]{2012ApJ...745...73N}. The relation between $v_{\rm neb}$ and other properties such as the rate of evolution of the velocity of Si absorption lines (the \emph{velocity gradient}) or the colors during the photospheric phase of evolution \citep{maeda10, 2011MNRAS.413.3075M, 2011A&A...534L..15C} seems to lend support to this picture. For the nearby SN 2011fe, there is evidence for lack of interaction between the ejecta and a companion star or a companion star wind. The radius of the companion star has been constrained to be less than 0.1 $R_{\odot}$ \citep{2012ApJ...744L..17B, 2011Natur.480..344N, 2011arXiv1110.2538B}, and its mass loss, less than $\approx 10^{-8}$ $M_{\odot}$ yr$^{-1}$ \citep{2012ApJ...746...21H}. However, \citet{2011Sci...333..856S} showed that SNe Ia show an excess of blue-shifted narrow Na lines, which is evidence for CSM around type Ia progenitors ejected before explosion. Particular cases, specifically SN 2002ic \citep{2003Natur.424..651H}, SN 2006X \citep[][c.f. \citealt{2008AstL...34..389C}]{2007Sci...317..924P}, SN 1999cl \citep{2009ApJ...693..207B} and SN 2007le \citep{2009ApJ...702.1157S} have also provided evidence for CSM around SNe Ia, although in the case of SN 2002ic the nature of the explosion is unclear \citep{2006ApJ...653L.129B}. \citet{2009ApJ...702.1157S} noted that SNe Ia with variable narrow absorption lines tend to have broad lines and high velocity gradients. \citet{2012ApJ...748..127F} recently suggested that higher velocity at peak, redder SNe Ia are preferentially found in lower mass host galaxies, but also that they have an excess of blue-shifted absorption systems \citep{2012arXiv1203.2916F}. These observations highlight the importance of understanding the relation between progenitor systems and dust properties for cosmological distance determinations \citep[e.g.][]{2008ApJ...686L.103G, 2010AJ....139..120F}. Since SNe Ia with $v_{\rm neb} \ge 0$ are known to have higher velocity gradients, generally associated to higher velocities at peak, and to be redder \citep{maeda10, 2011MNRAS.413.3075M}, we investigate the relation between the presence or absence of narrow Na I D1 \& D2 and Ca II H \& K absorption lines (hereafter Na and Ca lines), a proxy for material in the line of sight, with $v_{\rm neb}$, a geometrical proxy that should depend only on viewing angle \citep{maeda10} and that should therefore be independent of any evolutionary or host galaxy effects. | Our main conclusion is that the distribution of EWs from narrow blended Na and Ca lines differs significantly between SNe Ia with negative and positive $v_{\rm neb}$. Because the sign of $v_{\rm neb}$ should be purely geometrical, it should not correlate with any host galaxy or average SN properties. This suggests that part of the lines are formed by CSM ejected by the progenitor before explosion, which is found asymmetrically distributed after explosion. The differences in color found for SNe Ia with different $v_{\rm neb}$ \citep{2011MNRAS.413.3075M}, which we confirm using SNooPy and SiFTO light curve fitting, should be explained in part by extinction due to an asymmetric distribution of CSM, and not by different ejecta temperatures with viewing angle alone. There are two possibilities to explain this result: 1) the CSM material is asymmetrically ejected from the system \emph{before} explosion and is aligned with the side of the ejecta where ignition occurs, or 2) the CSM material is initially spherically symmetric and is affected by an asymmetric distortion \emph{after} explosion, which is stronger on the side where ignition occurs. One possible physical process that could break the symmetry of the WD and have an effect on both the direction of ignition and the distribution of CSM is rotation. However, rotation should produce cylindrically symmetric systems, where positive or negative $v_{\rm neb}$ should lead to the same average CSM properties in the line of sight. Another possibility is the gravitational field of a companion star, but the typical velocity offsets found for the absorbing material in \citet{2011Sci...333..856S} implies that the material was ejected at a time before explosion much longer than the expected orbital period of the progenitor systems, again producing some cylindrical symmetry in average. Thus, we conclude that only 2) is possible to explain our results. One way to produce an asymmetric distribution of dust, EW(Na) and EW(Ca) is to have an asymmetric ionizing field \emph{after} explosion. This would destroy most of the dust and ionize Na and Ca in the direction where ignition occurs. A comparison with pre--explosion images of the core--collapse SN 2012aw \citep{2012arXiv1204.1523F} supports the idea of dramatic changes in extinction after explosion. The fact that the correlations are found to be stronger for dust (colors) than for EW(Na) and EW(Ca) is suggestive, since they are in ascending order of ionization potential. One problem with this interpretation could be the constraints in radio and X--ray for CSM--ejecta interactions. This would be solved if the material were placed far enough from the ejecta, and only in some cases allow the ejecta to interact with the CSM, as was seen in SN 2002ic, SN 2006X and 2007le, or if the measured EWs are due to saturated lines in low mass, clumpy CSM with a wide distribution of velocities. Our results provide hints about the explosion asymmetry and CSM properties of SN Ia progenitors. If a larger sample confirms that most SNe Ia with $v_{\rm neb} \ge 0$ have high EWs of Na and Ca absorption lines, the majority of progenitors should contain significant CSM. If this is the case, most SNe Ia with $v_{\rm neb} < 0$ should have more strongly ionizing fields (higher energy photons) after explosion, either during shock breakout, or afterwards if more radioactive material is exposed in the ignition side of the explosion. Finally, significant CSM is consistent with some $M_{\rm Ch}$--SD or slow DD scenarios \citep[e.g.][]{2010ApJ...725..296F} and not with violent DD scenarios. However, in our interpretation it is difficult to explain the SNe Ia with $v_{\rm neb} \ge 0$ and low EWs, which may suggest that multiple progenitor scenarios are at work. | 12 | 6 | 1206.4697 |
1206 | 1206.3412_arXiv.txt | {We present the results of a visual search for galaxy-scale gravitational lenses in $\sim$7\degr$^2$ of Hubble Space Telescope (HST) images. The dataset comprises the whole imaging data ever taken with the Advanced Camera for Surveys (ACS) in the filter F814W (I-band) up to August 31$^{\rm st}$, 2011, i.e. 6.03\degr$^2$ excluding the field of the {\it Cosmic Evolution Survey} (COSMOS). In addition, we have searched for lenses in the whole Wide Field Camera 3 (WFC3) near-IR imaging dataset in all filters (1.01\degr$^2$ ) up to the same date. Our primary goal is to provide a sample of lenses with a broad range of different morphologies and lens-source brightness contrast in order to design and train future automated lens finders in view of all-sky surveys. Our criteria to select lenses are purely morphological as we do not use any color or redshift information. The final candidate selection is very conservative hence leading to a nearly pure but incomplete sample. We find 49 new lens candidates: 40 in the ACS images and 9 in the WFC3 images. Out of these, 16 candidates are secure lenses owe to their highly recognizable morphology, 21 more are very good candidates, and 12 more have morphologies compatible with gravitational lensing. The imaging dataset is heterogeneous in depth and spans a broad range of galactic latitudes. It is therefore insensitive to cosmic variance and allows to estimate the number of galaxy-scale strong lenses on the sky for a putative survey depth. Because of the incompleteness of the sample, the estimated lensing rates should be taken as lower limits. Using these, we anticipate that a 15 000\degr$^2$ space survey such as Euclid will find at least 60~000 galaxy-scale strong lenses down to a limiting AB magnitude of I=24.5 (10-$\sigma$) or I=25.8 (3-$\sigma$). } | Gravitational lensing in its weak and strong regimes is currently one of the best tools to study dark matter and dark energy \citep[e.g.][]{Hu1999}. It is also the most reliable way to weight precisely galaxies up to several effective radius \citep[e.g.][]{Gavazzi2007}. In combination with stellar dynamics, it provides the complementary measurement to determine the lens mass profile \citep[e.g.][]{Barnabe2009} and to break the disc-halo degeneracy in spiral galaxies \citep[e.g.][]{Dutton2011}. Thanks to large sample of strong lenses, statistical studies of galaxy mass properties and evolution with redshift are feasible \citep[e.g.][]{faure2011,faure2009,lagattuta2010,Auger2010, Koopmans2006}. In some cases, when the radial extent of an Einstein ring is particularly large or when sources at multiple redshifts are lensed by the same object \citep[e.g.][]{Gavazzi2008}, the measurement of the mass slope in the lens can be of exquisite quality and allow to constrain the cosmological parameters \citep[e.g.][]{Suyu2010} in a way fully competitive with other cosmological probes \citep[see][for a short comparison of the methods]{Suyu2012}. Even with a few hundreds of strong lenses in hand, current samples remain small and must be drastically extended to allow one for high precision statistical studies and to provide a genuine understanding of the distribution of visible and dark matter in galaxies. The increasing number of optical all-sky surveys either from the ground or from space, lends considerable hope to build such large samples, with tens or even hundreds of thousands of objects. However, the automated techniques available to find the lenses from petabytes of imaging data are so far limited in their efficiency and tend to produce a large number of false positives that require significant post-processing cleaning. Among the best automated robots to find lenses are "Arcfinder" \citep{Seidel2007}, which was primarily developed to find large arcs behind clusters and groups, and the algorithm by \citet{Alard2006} used by \citet{Cabanac2007} and \citet{More2012} which was optimized to look for arcs produced by individual galaxies and groups in the CFHT Strong Lensing Legacy Survey (SL2S). While the latter techniques do not rely on any model for the mass distribution in lens galaxies, other automated robots consider any galaxy as a potential lens and predict where lensed images of a background source should be before trying to identify them on the real data \citep{Marshall2009}. These robots are being built and trained to mimic the performances of the human eye/brain in recognizing a lens system among other astronomical objects. To perform this training efficiently, we need to build strong lens samples with large varieties of image configurations. % So far, the most efficient way in this respect is the visual inspection of a significant portion of the sky, with the goal of providing a test bench to devise, test and calibrate automated robots. Such robots will be mandatory to carry out lens searches in surveys of intractable size for a few astronomer , i.e., the all-sky surveys that may take place in the next decade. An alternative to robots will be the search for strong lenses by citizens in projects like the Galaxy Zoo \citep{lintott2008} and the future Lens Zoo (PI: P. Mashall). But we forecast that their contribution will be complementary to that of robots, in the action of classifying the robots numerous outputs. \begin{figure*}[ht!] \includegraphics[width= 1.0\textwidth]{./all_sky_data.jpeg} \caption{Distribution of the ACS (left) and WFC3 (right) images on the sky. The red dots indicate the position of the fields with a lens candidate. On the left panel, each point corresponds to a single ACS field of 11\arcmin$^2$. The total area covered by the ACS is 6.03\degr$^2$. On the right panel, each point corresponds to a single WFC3 field of 4.65\arcmin$^2$. The total area covered by the WFC3 survey is 1.01\degr$^2$. } \label{fig:allsky} \end{figure*} Visual inspections of HST archive images in search for strong lenses have been successfully attempted in the past: \citet{Ratnatunga1999} found 10 lenses in the HST medium deep survey, followed by \citet{Moustakas2007} who found another few systems. So far, the largest search for strong lenses in HST images is the one conducted in the 1.64\degr$^2$ field of view of the {\it Cosmic Evolution Survey} (COSMOS) \citep{scoville2007} by \citet{Faure2008} and \citet{Jackson2008}. They found in total 179 lens candidates among which 22 display multiple images of the source and/or have both lens and source redshifts to confirm their lens nature \citep[see][]{faure2011}. In the present paper, we describe a search for strong lenses by visual inspection in all the HST images ever taken with the Wide Field Channel (WFC) of the Advanced Camera for Surveys (ACS) through the F814W filter (I-band). The total field of view explored is 6.03\degr$^2$ and {\it excludes} the COSMOS field. Since future wide field surveys will include near-IR imaging, we also carry out the experiment using the data taken with the Wide Field Camera 3 (WFC3) in the near-IR channel and the F160W filter. As the WFC3 is a rather new instrument on the HST and since the field of view of the camera itself is limited, the total area we cover in the near-IR is smaller than with the ACS data, i.e., 1.01\degr$^2$. Our goal with this work is both to provide a training set of lenses spanning a broad diversity in image configuration, hence helping to design automated searches and to estimate empirically the number counts of strong lenses as a function of depth. Throughout the paper, all magnitudes are in the AB magnitude system. | We have conducted a visual search for galaxy-scale strong gravitational lenses in a total of 7 ${\rm deg}^2$ of the sky, using archival HST images. The data comprise the whole ACS/WFC F814W imaging data available as of August 31$^{\rm st}$, 2011 (6.03 ${\rm deg}^2$, excluding the COSMOS field) as well as the WFC3 near-IR images (1.01 ${\rm deg}^2$ in all filters) up to the same date. We have found 49 new lens candidates: 40 in the ACS images and 9 in WFC3 images. Out of these, 16 are without any doubt genuine lenses, owing to their striking morphology. An additional 21 are excellent candidates and 12 have morphologies compatible with strong lensing but would need further investigation before being confirmed. If these candidates are indeed genuine lenses, they will contribute to broaden the variety of known lens morphologies and may prove useful to devise an train automated lens finders for ongoing and future all-sky surveys. Our search for strong lenses is purely based on morphological criteria, with limited color information for 16 objects out of 49. It is not aimed at providing an homogeneous sample well suited for cosmology or galaxy evolution study, but rather to serve as a training set spanning a broad range of lensing configurations and useful for future automated searches. From our search, we conclude that, while the human-eye search may detect unusual lenses, such as faint arcs hidden in the glare of bright lenses or in complex/multiple lenses, it also has its specific biases as some genuine lenses will be discarded if only morphological criteria are used. The large amount of data in the optical F814W band allows us to estimate the lensing rate as a function of depth, at the spatial resolution of the HST. We believe that these counts are lower estimates of the true lens counts owing to the high purity of the sample and probable incompletness. Scaling our lensing rate, we estimate that a 15~000${\degr}^2$ optical survey such as Euclid will find at least 60~000 strong lenses down to I=25.8~mag (3-$\sigma$), which is about 400 times the current number of known strong lenses in the same area of sky. \begin{table*}[t!] \label{Table:1} \caption{ List of lens candidates found from the HST/ACS data. Column 1: Identity number. Column 2: Lens candidate name. Column 3 and 4: Coordinates. Column. 5: Radius of the arc in arc-second. Column 6: Total magnitude of the lensing galaxy in the F814W band, integrated in the 2-D light profile. The typical errors are 1 to 2$\%$. Column 7: Effective radius in arcsec. The typical errors are 1 to 2$\%$. Column 8: Lensing galaxy ellipticity, $\epsilon=($b/a$)$. The typical errors are 1 to 2$\%$. Column 9: Lensing galaxy position angle. Column 10: Comment, it describes the lensing candidate, see Section~\ref{class} for details. For candidate \#33 (ACS J153234.81+324115.3) no satisfactory fit was obtained. The unrealistic effective radius is not reported nor plotted in the Fig.~\ref{fig:radii}} \begin{tabular}{llllllllll} \hline \hline \vspace{0.5 mm}\\ ID&Name&RA&DEC&r$_{arc}$&I $($lens$)$&R$_{eff}$& $\epsilon$ &PA&Comment\\ & &h:m:s&$^{\circ}$:$\arcmin$:$\arcsec$ & $\arcsec$& F814W&$\arcsec$ & & &\\ \hline \vspace{0.5 mm}\\ 01&ACS J001423.02$-$302109.8&00:14:23.02&$-$30:21:09.8&1.52&18.86&0.25&0.70&$+$39.1&A1b\\ 02&ACS J001426.26$-$302255.9&00:14:26.26&$-$30:22:55.9&1.00&18.06&0.70&0.61&$+$30.2&A1a\\ 03&ACS J005656.81$-$273841.0&00:56:56.81&$-$27:38:41.0&0.93&20.34&1.03&0.29&$+$15.3&A1b\\ 04&ACS J010559.53$-$254438.0&01:05:59.53&$-$25:44:38.0&0.88&20.56&0.76&0.85&$+$41.5&A1b\\ 05&ACS J011018.22+193819.5&01:10:18.22&+19:38:19.5 &3.10&17.03&2.47&0.90 &$-$32.7&A1a\\ 06&ACS J021634.62$-$051035.3&02:16:34.62&$-$05:10:35.3&1.69&18.28&0.57&0.44&$+$06.1&A1b\\ 07&ACS J084710.65+344826.4&08:47:10.65&+34:48:26.4&1.98&20.25&0.66&0.88&$-$74.1&A1a\\ & & & & &22.13&0.04&0.34&$+$23.3\\ 08&ACS J093323.88+441549.7&09:33:23.88&+44:15:49.7&1.32&19.71&0.26&0.93&$-$80.1&A1b\\ 09&ACS J095139.44+684731.2&09:51:39.44&+68:47:31.2&0.77&19.78&0.50&0.71&$+$49.2&A1a\\ 10&ACS J095718.69+685731.5&09:57:18.69&+68:57:31.5&0.82&18.62&1.29&0.91&$+$85.8&A1b\\ 11&ACS J103751.40$-$124327.5&10:37:51.40&$-$12:43:27.5&1.48&19.25&2.04&0.54&$+$21.1&A1b\\ & & & & &19.99&0.68&0.54&$-$03.5\\ 12&ACS J104701.80+171739.7&10:47:01.80&+17:17:39.7& 0.46&19.34&0.86&0.72&$+$64.5&P2b\\ 13&ACS J111052.23+284233.7&11:10:52.23&+28:42:33.7&0.79&18.67&0.76&0.80&34.51&A2c\\ % 14&ACS J111127.06+285000.2&11:11:27.06&+28:50:00.2&0.39&21.02&1.50&0.69&$+$62.3&A1c\\ 15&ACS J111715.58+174500.0&11:17:15.58&+17:45:00.0&0.95&19.68&0.49&0.82&$+$49.1&P2b\\ 16&ACS J112310.02+141020.2&11:23:10.02&+14:10:20.2&1.01&20.50 &0.18&0.50&$-$07.7&A1c\\ 17&ACS J115135.05+004912.4&11:51:35.05&+00:49:12.4&1.87&19.10 &0.16&0.77&$+$08.2&A1c\\ 18&ACS J115138.59+004848.4&11:51:38.59&+00:48:48.4&1.96&19.12&1.24&0.41&$-$04.3&A1c\\ 19&ACS J121637.27$-$120042.2&12:16:37.27&$-$12:00:42.2&1.46&21.27&0.16&0.77&$-$38.5&A1b\\ 20&ACS J122332.64$-$123940.3&12:23:32.64&$-$12:39:40.3&0.36&19.68&0.19&0.54&$+$53.2&A1a\\ 21&ACS J130042.73+280523.3&13:00:42.73&+28:05:23.3&1.08&17.56&1.46&0.62&$-$32.1&A1b\\ 22&ACS J135401.12$-$123450.4&13:54:01.12&$-$12:34:50.4&2.95&19.32&1.35&0.83&$-$2.80&A1c\\ 23&ACS J140237.11+542716.4&14:02:37.11&+54:27:16.4&1.85&17.99&1.38&0.71&$-$25.0&A1a\\ 24&ACS J140339.94+541633.3&14:03:39.94&+54:16:33.3&0.90&18.66&3.44&0.34&$+$40.9&A1a\\ 25&ACS J141649.81+522951.7&14:16:49.81&+52:29:51.7&0.34&20.82&0.94&0.96&$+$06.9&A1b\\ 26&ACS J141657.72+433026.3&14:16:57.72&+43:30:26.3&0.78&20.82&0.16&0.79&$-$74.5&A1b\\ 27&ACS J141710.17+433133.2&14:17:10.17&+43:31:33.2&2.14&18.16&4.94&0.39&$+$20.7&A1a\\ 28&ACS J142028.48+525541.5&14:20:28.48&+52:55:41.5&0.61&19.67&0.56&0.91&$-$77.8&A1b\\ 29&ACS J142158.70+531045.9&14:21:58.70&+53:10:45.9&0.82&22.38&0.40&0.59&$+$41.4&A1c\\ 30&ACS J143542.20+342841.9&14:35:42.20&+34:28:41.9&0.96&19.96&1.77&0.43&$-$62.4&A1c\\ 31&ACS J144000.07+313123.5&14:40:00.07&+31:31:23.5&0.60&22.08&0.12&0.61&$+$30.6&A1b\\ 32&ACS J152702.08+355726.8&15:27:02.08&+35:57:26.8&1.73&19.59&1.16&0.68&$+$84.5&A1c\\ 33&ACS J153234.81+324115.3&15:32:34.81&+32:41:15.3&1.21&18.54&16.67&0.38&$+$31.6&P4b\\ 34&ACS J171750.68+593457.5&17:17:50.68&+59:34:57.5&1.79&17.92&0.68&0.77&$+$09.0&A1b\\ 35&ACS J171817.43+593146.4&17:18:17.43&+59:31:46.4&2.35&15.21&2.72&0.79&$+$58.5&A1a\\ 36&ACS J173923.28+690706.6&17:39:23.28&+69:07:06.6&1.23&20.94&0.14&0.38&$+$78.5&A1b\\ 37&ACS J174035.58+690328.1&17:40:35.58&+69:03:28.1&1.12&19.08&0.63&0.31&$+$79.2&A1a\\ 38&ACS J221501.12$-$135822.9&22:15:01.12&$-$13:58:22.9&0.82&19.10 &0.78&0.22&$-$57.5&A1a\\ 39&ACS J230325.68+085212.6&23:03:25.68&+08:52:12.6&0.72&17.62&1.08&0.35&$-$01.0&A1a\\ 40&ACS J235130.60$-$261459.7&23:51:30.60&$-$26:14:59.7&3.32&16.26&2.49&0.74&$-$38.0&A1a\\ \hline \end{tabular} \label{tab:acscan} \end{table*} \begin{table*}[h!] \label{Table:2} \caption[]{List of lens candidates found from HST/WFC3 data. The description of the columns is given in the caption of Table~\ref{tab:acscan} except that the magnitudes are for the filter given in the "Filter" column. \label{tab:WFC3can} } \begin{tabular}{lllllllllll} \hline \hline \vspace{0.5 mm}\\ ID&Name&RA&DEC&r$_{arc}$&R$_{eff}$& $\epsilon$ &PA&Mag&Filter&Comment\\ & &h:m:s&$^{\circ}$:$\arcmin$:$\arcsec$ &$\arcsec$ &$\arcsec$&& &\\ \hline \vspace{0.5 mm}\\ 1&WFC3 J023924.56$-$013600.82&02:39:24.56&$-$01:36:00.82&0.91&0.65&0.60&$-$56.5&20.50&F140W&R1b\\ 2&WFC3 J033227.44$-$275520.42&03:32:27.44&$-$27:55:20.42&1.83&1.85&0.57&$+$75.9&18.35&F105W&A1a\\ 3&WFC3 J033245.17$-$274940.46&03:32:45.17&$-$27:49:40.46&1.71&0.25&0.11&$+$67.9&20.48&F125W&A1a\\ 4&WFC3 J100023.53+021653.34&10:00:23.53&+02:16:53.34&2.08&1.29&0.50&$-$49.8&19.33&F140W&A1b\\ 5&WFC3 J140221.89+094520.16&14:02:21.89&+09:45:20.16&1.04&3.67&0.78&$-$73.0&20.13&F110W&A2c\\ 6&WFC3 J142030.50+530249.23&14:20:30.50&+53:02:49.23&0.73&0.35&0.57&$-$81.4&22.84&F125W&A1c\\ 7&WFC3 J143703.21+350153.55&14:37:03.21&+35:01:53.55&0.69&0.55&0.64&$+$40.7&19.87&F160W&R1a\\ 8&WFC3 J171336.64+585640.73&17:13:36.64&+58:56:40.73&1.43&5.98&0.87&$+$45.0&19.07&F160W&A1c\\ 9&WFC3 J171736.34+601437.60&17:17:36.34&+60:14:37.60&2.23&0.23&0.69&$-$12.6&20.69&F160W&A1b\\ \hline \end{tabular} \end{table*} \begin{table}[h] \label{Table:3} \caption{Image filters used to create the color images in Fig.~\ref{fig:ACScolor}. When only two filters are available, the mean of the two images was taken as the green channel.} \label{tab:colorfilters} \begin{tabular}{llll} \hline \hline Name& Red & Green & Blue\\ \hline ACS J001423.02$-$302109.8&F814W & F606W & F435W\\ ACS J001426.26$-$302255.9&F814W & F606W& F435W\\ ACS J021634.62$-$051035.3&F814W&\_&F606W\\ ACS J095139.44+684731.2& F814W&\_& F606W\\ ACS J104701.80+171739.7& F814W&\_& F435W\\ ACS J111052.23+284233.7& F814W&\_&F555W\\ % ACS J115135.05+004912.4& F814W&\_&F435W\\ ACS J130042.73+280523.3& F814W&\_&F475W\\ ACS J140237.11+542716.4& F814W& F555W& F435W\\ ACS J140339.94+541633.3& F814W& F555W& F435W\\ ACS J141649.81+522951.7& F814W&\_&F606W\\ ACS J141657.72+433026.3& F814W&\_&F606W\\ ACS J141710.17+433133.2& F814W&\_& F606W\\ ACS J142028.48+525541.5& F814W&\_&F606W\\ ACS J142158.70+531045.9& F814W&\_& F606W\\ ACS J144000.07+313123.5& F814W&\_& F606W\\ ACS J230325.68+085212.6& F814W&\_&F435W\\ \hline \end{tabular} \end{table} | 12 | 6 | 1206.3412 |
1206 | 1206.5002_arXiv.txt | The discovery of multiple evolutionary sequences has challenged the paradigm that globular clusters (GCs) host simple stellar populations. In addition, spectroscopic studies of GCs show a spread in light-element abundances, suggesting that multiple sequences can be formed from gaseous ejecta processed in evolved cluster stars. If multiple sequences originate from within GCs, then it should be determined how such stellar systems retain gas, form new stars within them and subsequently evolve. Here we expand upon previous studies and carry out hydrodynamical simulations that explore a wide range of cluster masses, compactness, metallicities and stellar age combinations in order to determine the ideal conditions for gas retention. We find that up to 6.4\% of the mass of the star cluster can be made up of retained stellar wind gas at the time star formation is triggered. However, we show that multiple episodes of star formation can take place during the lifetime of a star cluster in particular for times $\gtrsim 1$~Gyr, thus leading to a sizable enhancement in the total number of new stars. The fact that this favorable star formation time interval coincides with the asymptotic giant branch (AGB) phase seems to give further credence to the idea that, at least in some GCs, there are stars which have formed from material processed by a previous generation of stars. The ability of extended heating sources, such as pulsar outflows or accretion onto compact objects, to hamper gas retention is illustrated via a simple numerical treatment. | The idea that stars in globular clusters (GCs) are a coeval population has been a long disputed topic with significant discussion devoted to the presence (or lack thereof) of multiple main sequences and subgiant branches \citep[see][for a review]{piotto2009}. However, recent observational evidence has demonstrated that these multiple stellar populations (MSPs) are not only ubiquitous but also make up a sizable fraction ($\sim$40-60\%) of the total stellar mass \citep{caloi2007,dantona2008}. Spectroscopic observations of MSPs in GCs have also unveiled differences in their light element abundances \citep[see e.g.][]{gratton2004,deSilva2009,martell2009}. These anomalies can only be produced if these MSPs were formed from material processed at temperatures $T> 10^7$K \citep{dercole2008,renzini1981,ventura2008b}. To explain the presence of MSPs with the observed abundance anomalies the following timeline has been identified \citep{conroy2011,cottrell1981,smith1987,carretta2010b}: after the first generation of stars is formed and intracluster gas is expelled by supernovae, matter processed in the interiors of AGB stars is returned to the ISM through winds and after several 100 Myrs a second generation of stars is formed from a mixture of stellar-processed material and captured ISM gas. Some of the issues at the forefront of attention include the type and age of stars able to process gas at the required internal temperatures while at the same time efficiently returning material into GCs to form a sizable mass fraction of new stars. We address both of these issues here. The effective retention of stellar wind material to create MSPs of comparable mass presents a challenge to the current understanding of GC formation \citep{conroy2011b}. Accounting for the inflow and mixing of pristine ISM material slightly lowers the retained mass requirements, however the original GCs still need to be significantly more massive than those currently observed \citep{naiman2011,pflamm2009,dercole2010,conroy2011}. Despite their relevance, stellar wind gas retention has so far only been studied for a very restricted number of systems \citep{dercole2008,dercole2010,conroy2011b,conroy2011}. Here we expand upon these studies with hydrodynamical simulations which explore the possible combinations of cluster mass, compactness and stellar age to determine the optimum parameter space for maximal gas retention. | \label{section:discussion} Throughout this work, we have used simplified models to determine how effective a star cluster of a particular age is at retaining gas emanating from its stellar members, under the assumption that its potential remains unaltered during the simulated phase. We have further assumed a single metallicity for all clusters and have disregarded any heating sources besides supernova feedback and the stellar winds themselves. In this section we relax both of these assumptions. \subsection{Metallicity} Cluster to cluster variations in light element abundances are commonly observed \citep{caldwell,beasley}. These variations may cause changes in the mass loss histories of the individual stellar members and the cooling properties of the shocked gas. As the stellar mass loss prescription are mostly independent of metallicity during the evolutionary time periods that are conducive to star formation ($0.1 \, {\rm Gyrs} \lesssim t_{i} \lesssim 100 \, {\rm Gyrs}$) \citep{blocker}, we account for the effects of varying metallicity solely in the cooling function. Figure \ref{fig:fig4} shows the effects of changing metallicity for a typical star forming cluster, described here by $M_c = 10^7 \, M_\odot$, $\sigma_v = 50 \, {\rm km/s}$, and $t_{i} = 212 \, {\rm Myrs}$. As the metallicity is decreased, the cooling becomes less efficient and more material is allowed to flow into the center of the cluster before catastrophic cooling occurs. This enables relatively more massive star forming episodes to be triggered. Interestingly, when metallicity is decreased beyond $Z < 10^{-1}$, the cooling becomes weak enough to prevent catastrophic cooling at times $\leq t_{i}$. These results suggest that the range of cluster parameters over which large central densities will persist before catastrophic cooling takes place (Figure \ref{fig:fig4}) will depend on the metallicity of the emanating stellar winds, though, as illustrated in Figure \ref{fig:fig7}, the differences are not marked. \subsection{Intercluster Heating Sources} In addition to altering the cooling curves, the inclusion of additional cluster heating sources may prevent effective gas retention in our simulations. We address this problem here by artificially increasing the energy input rate: $q_{\epsilon,{\rm new}} = (1+H)q_\epsilon = (1+H)\frac{1}{2} q_m(r)v_w^2$. Under this assumption, the additional heating sources follow the potential's stellar distribution. Figure \ref{fig:heating} shows the effects of the additional heat input in one of our otherwise star forming simulations. For $H < 2.0$, the gas in the simulation still collapses, triggering star formation. For larger values, on the other hand, the cluster is unable to effectively retain the gas and, as a result, star formation never ensues. By integrating $q_{\epsilon,{\rm new}}$ over the cluster's core for $H=2.0$, we derive the total energy input rate required to overturn the central mass build up, which for this simulation is about $10^{35} \, {\rm erg s^{-1}}$. In many cases, the additional energy injection sources might not follow the stellar distribution. As an example, let's compare the heat distribution expected from accreting neutron stars under the assumption that the accretion feedback is proportional to the Bondi accretion rate: $q_{\epsilon,{\rm ns}} \propto \dot{M} \propto \rho(r) T(r)^{-3/2}$ \citep{bondi}. Using the volume-averaged density and temperature in the cluster core, $\bar{\rho} \approx 10^{-22} \, {\rm g\; cm^{-3}}$ and $\bar{T} \approx 10^4$K, we derive the average luminosity of a single, accreting neutron star: $L_{NS} \approx 10^{33} \, {\rm ergs\; s^{-1}}$. This implies that $\gtrsim$ 100 accreting neutron stars are required to reside in the cluster's core in order to significantly offset its cooling properties. However, to accurately test this phenomena a multi-dimensional approach would be required as feedback would not necessarily act as a simple heating prescription. In this work, we have also tried to minimize the effect of external mass inflow by considering star clusters in isolation. This is certainly not the case for clusters moving through cold, dense environments, as they can potentially amass a significant amount of gas from their surroundings \citep{naiman2009,naiman2011}. The stellar winds and exterior inflows in such clusters could combine to create even larger central density enhancements \citep{naiman2009,naiman2011,conroy2011,priestley2011,pflamm2009}. A self consistent treatment of both interior and exterior gas accumulation in multi-dimensional simulations, which includes the effects of compact object accretion, tidal stripping, photoionization, and pulsar heating will be presented elsewhere. While previous work has estimated the ability of star clusters to retain stellar winds \citep{dercole2008,dercole2010,conroy2011b,conroy2011}, calculations have so far been restricted to a small range of cluster properties and stellar ages. Motivated by this, we have calculated gas retention in star clusters of various ages, stellar mass and compactness. In agreement with previous studies we find that before star formation is triggered about $\leq$10\% of the total cluster mass is comprised of retained stellar wind gas, naively implying that the original proto-GCs had to be more massive than what is observed today \citep{conroy2011b,dercole2008}. However, we show that multiple episodes of star formation can in fact take place during the lifetime of a cluster in particular between $\sim 1$ and $\sim 10$~Gyrs, thus suggesting a sizable enhancement in star formation. The overlap of this time range with the AGB phase further strengthens the case for stellar wind retention as a critical component in the formation of subsequent generations of stars in GCs. We thank Mark Krumholz, Charlie Conroy, David Pooley, Rebecca Bernstein, Chung-Pei Ma and Morgan Macleod for useful discussions. The software used in this work was in part developed by the DOE-supported ASCI/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago. Computations were performed on the Pleaides UCSC computer cluster. This work is supported by NSF: AST-0847563 (J.N. and E.R.), NASA: NNX08AL41G (J.N. and D.L.), and The David and Lucile Packard Foundation (E.R.). \begin{figure*} \centering\includegraphics[width=0.6\textwidth]{fig1.eps} \caption{Average cluster mass loss rates and wind velocities as a function of time for $Z = 1/10 Z_\odot$. Green lines assume gas dynamics are dominated by the wind properties of the turn off stars, black lines show the population averaged values. Three representative times in the clusters age are denoted by the solid vertical lines. The current age of M15 is denoted by the vertical dashed line.} \label{fig:fig1} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.95\textwidth, angle=270]{fig2_t3.eps} \caption{Hydrodynamic profiles for the three representative times denoted in Figure \ref{fig:fig1}. These three plots are the density, temperature and velocity profiles at three representative times for a $M_c = 10^7 \, M_\odot$, $\sigma_v = 30 \, {\rm km/s}$ model. Solid and dashed lines represent the turn off mass and population averaged prescriptions, respectively. } \label{fig:fig2} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.95\textwidth, angle=270]{fig3.eps} \caption{Hydrodynamic profiles for M15 with $M_c = 4.5 \times 10^5 \, M_\odot$ and $\sigma_v = 14 \, km/s$ \citep{mcnamara,gerssen} at the three representative times in Figure \ref{fig:fig1}, including the predicted profiles for its current age (dashed lines). Here, the population averaged prescription is used.} \label{fig:fig3} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.99\textwidth, angle=270]{fig4.eps} \caption{Mass accumulation, $M_{\rm acc}$, as a function of the potential parameters $M_{\rm c}$ and $\sigma$ for the turn off mass prescription. The blue shaded regions delimit the boundaries for which our models collapse and trigger star formation. {\textit{Right:}} Contours of $M_{\rm acc}/M_{\rm c}$ for a fixed cluster core mass of $M_{\rm c} = 10^7 \, M_\odot$ and a varying velocity dispersion. {\textit{Left:}} Contours of $M_{\rm acc}/M_{\rm c}$ for a fixed velocity dispersion of $\sigma_v = 27 \, {\rm km/s}$ and varying cluster mass. } \label{fig:fig4} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.99\textwidth, angle=270]{fig5.eps} \caption{Mass accumulation, $M_{\rm acc}$, as a function of $\sigma$ for a fixed cluster core mass of $M_{\rm c} = 10^7 \, M_\odot$, calculated using the population averaged prescription. The blue shaded region delimits the boundaries for which our models collapse and trigger star formation.} \label{fig:fig5} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.99\textwidth,angle=270]{snExplosionsSmall.eps} \caption{Cycles of mass accumulation in a potential with $M_{\rm c} = 10^7 \, M_\odot$, $\sigma_v = 26 \, {\rm km/s}$ and $t_i = 2000 \, {\rm Myrs}$ including star formation, and supernovae explosions. Mass accumulates till star formation is triggered, then $90 \%$ of this mass is assumed to form stars, removing the gas from the simulation (dashed lines). Supernovae persist for 100~Myrs at a rate of $2 \times 10^{-6} \, {\rm yr^{-1}}$ further stripping mass from the system. For simplicity, we have assumed the mass injection properties to be steady during the simulation.} \label{fig:fig6} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.99\textwidth,angle=270]{fig7.eps} \caption{Gas properties for $M_c = 10^7 \, M_\odot$, $\sigma_v = 30 \, {\rm km/s}$ cluster at $t_{i} = 313 \, {\rm Myrs}$ for different metallicities. Here, the population averaged values of the wind parameters are used.} \label{fig:fig7} \end{figure*} \begin{figure*} \centering\includegraphics[height=0.99\textwidth, angle=270]{heatingPlotOut.eps} \caption{{\textit{Left 3 Plots:}} Gas properties for $M_c = 10^7 \, M_\odot$, $\sigma_v = 30 \, {\rm km/s}$ cluster at $t_{i} = 1860 \, {\rm Myrs}$ with added heat in the form of thermal energy: $e_{int,new} = (1 + H)e_{int,old}$. Here, the population averaged values of the wind parameters are used.} \label{fig:heating} \end{figure*} | 12 | 6 | 1206.5002 |
1206 | 1206.0557_arXiv.txt | {Searching for the non-Trojan Jupiter co-orbitals we have numerically integrated orbits of 3\,160 asteroids and 24 comets discovered by October 2010 and situated within and close to the planet co-orbital region. Using this sample we have been able to select eight asteroids and three comets and have analyzed their orbital behavior in a great detail. Among them we have identified five new Jupiter co-orbitals: \cu, \sa, \ql, \gh, and \Larsen, as well as we have analyzed six previously identified co-orbitals: \hr, \ug, \qq, \aee, \wc\ and \ar. \cu\ is currently on a quasi-satellite orbit with repeatable transitions into the tadpole state. Similar behavior shows \gh\ which additionally librates in a compound tadpole-quasi-satellite orbit. \ql\ and \Larsen\ are the co-orbitals of Jupiter which are temporarily moving in a horseshoe orbit occasionally interrupted by a quasi-satellite behavior. \sa\ is moving in a pure horseshoe orbit. Orbits of the latter three objects are unstable and according to our calculations, these objects will leave the horseshoe state in a few hundred years. Two asteroids, \qq\ and \aee, are long-lived quasi-satellites of Jupiter. They will remain in this state for a few thousand years at least. The comets \ar\ and \wc\ are also quasi-satellites of Jupiter. However, the non-gravitational effects may be significant in the motion of these comets. We have shown that \wcs\ is moving in a quasi-satellite orbit and will stay in this regime to at least 2500 year. Asteroid \hr\ will be temporarily captured in a quasi-satellite orbit near 2050 and we have identified another one object which shows similar behavior - the asteroid \ug, although, its guiding center encloses the origin, it is not a quasi-satellite. The orbits of these two objects can be accurately calculated for a few hundred years forward and backward.}{Minor planets, asteroids: general - Comets: general - Celestial mechanics - Methods: numerical} | Small bodies locked in a co-orbital region are usually associated with the Lagrangian equilibrium points. For more than one century the Lagrange points were only a subject of theoretical considerations. The first object, Trojan asteroid (588) Achilles, that traces tadpole-shaped trajectory (TP) around Jupiter L$_4$ lagrangian point was discovered by Max Wolf in 1906. Trojan asteroids have also been discovered for Mars (Connors et al. 2005) and Neptune (Sheppard and Trujillo 2006). Shortly after the discovery of the first Trojan, Brown (1911) indicated, that stable horseshoe (HS) orbits, that surround L$_4$, L$_5$ and L$_3$ Lagrangian points can exist. Another kind of 1:1 resonant trajectories, not associated with the Lagrangian points, recently known as quasi-satellite (QS) orbits (Lidov and Vashkov'yak 1994, Mikkola and Innanen 1997), were predicted by Jackson (1913). All these families of orbits can be simply classified through the librational properties of the principal resonant angle, $\sigma$, which is defined by $\sigma=\lambda-\lambda_p$, where $\lambda$ and $\lambda_p$ are the mean longitude of the object and the planet respectively. TP orbits librate around $\pm 60^{\circ}$, but for eccentric TP orbits their libration centers are displaced with respect to the equilateral locations at $\pm 60^{\circ}$ (Namouni and Murray 2000). HS orbits librate around $180^{\circ}$ whereas QS librate around $0^{\circ}$. In the planar case, the families of TP, HS and QS orbits are disjoint, however, in a case of sufficiently inclined and eccentric orbits there can exist compound orbits which are unions of the HS (or TP) and QS orbits, and transient co-orbital orbits for which transitions between different families occur (Namouni 1999, Namouni et al. 1999, Christou 2000, Brasser et al. 2004a, Wajer 2009, 2010). This type of resonant behavior has been observed in the motion of several known near-Earth asteroids. We searched numerically for objects which are, or will be in the near future (in the next 100 years), co-orbitals of Jupiter and experience transitions between different types of co-orbital motion. We have selected these objects which have astrometric data collected over a longer period than 1 year and we have identified among them seven asteroids and one comet. Additionally, we decided to investigate three more objects of a shorter interval of observations (asteroid \aee\ and two comets: \ar\ and \wc\ which were previously investigated by Kinoshita and Nakai (2007). | \label{summary} We have analyzed the orbits of six already known, and we have identified five new non-Trojan, co-orbitals of Jupiter. These object have been divided into several classes according to their dynamical properties. We selected four such classes of the co-orbital behavior of the objects: TP$/$QS, HS$/$QS, long lasting QS and temporary co-orbitals. The most important results of our investigation are summarized in the Table 2. In this table are listed the Tisserand parameter with respect to Jupiter, the type of co-orbital motion, time of integration and the estimated period of predictability of the all considered objects. The analyzed objects can also be classified according to their Tisserand parameter with respect to Jupiter, $T_J$. According to this criterion cometary orbits are defined as those having $T_J < 3$, while asteroidal orbits are those with $T_J > 3$. Therefore, all the objects with $T_J < 3$ that do not present any signature of cometary activity are defined as an asteroids in cometary orbits (ACO) (Licandro et al. 2006). They also found that ACO with $T_J>2.9$ have spectra typical of the main belt objects while those with $T_J < 2.9$ shown comet-like spectra. All objects investigated here have $T_J < 2.9$, with the largest value of $T_J = 2.89$ for asteroid 2006 SA$_{387}$. The three analyzed comets belong to the Jupiter family comets ($T_J>2.0$) and, as one can see in the Table 2, the Tisserand parameter of analyzed asteroids also locate all of them as Jupiter family comets. In fact, the comets \wc\ and \ar\ were discovered as asteroids 2003 WC$_7$ and 2002 AR$_2$, and later a cometary activity was recognized on these objects (Kinoshita and Nakai 2007). One may expects that the analyzed asteroids, in the future, also can exhibit cometary outgassing. \begin{table*} \tiny \caption{\label{tab:dynamics}The Tisserand parameter, type of co-orbital behavior, integration and predictability period for the analyzed objects.} \begin{tabular}{l*{4}{c}r} Object & T$_J$ & Dynamical & Integration & \multicolumn{2}{c}{Predictability period} \\ & & behavior & period & from & to \\ \hline (241944) 2002 CU$_{147}$ & 2.60 & transient TP-QS & 1000-7000 & $<$1000&6500 \\ \hline 2007 GH$_{6}$ & 2.63 & transient TP-QS, compound TP-QS-TP & 1000-7000 & $<$1000&5500\\ \hline 2006 QL$_{39}$ & 2.53 & temporary compound HS & 1000-3000 & 1200&2791 \\ \hline 2006 SA$_{387}$ & 2.89 & temporary HS & 1000-3000 & 1780&2250 \\ \hline 2001 QQ$_{199}$ & 2.37 & long-lasting QS & 0-12000 & $<$0&$>$12000 \\ \hline 2004 AE$_{9}$ & 2.50 & long-lasting QS & 0-12000 & $<$0&$>$12000 \\ \hline (118624) 2000 HR$_{24}$ & 2.80 & temporary QS & 1000-3000 & 1350&2350 \\ \hline 2006 UG$_{185}$ & 2.72 & temporary co-orbital & 1000-3000 & 1800&2320 \\ \hline\hline 200P Larsen & 2.74 &transient QS-HS & 1000-3000 & 1917&2800 \\ \hline C/2002 AR$_{2}$ LINEAR & 2.52 & long-lasting QS & 1000-3000 & $<$1000& $>$3000 \\ \hline C/2003 WC$_{7}$ LINEAR-Catalina & 2.36 & compound TP-QS$/$QS & 1000-3000 & 1300&2500 \\\hline \end{tabular} \end{table*} The predictability period mainly depends on the close approaches with the Jupiter. As we can see in table 2 the most predictable orbits have objects which avoid close encounters with this planet, i.e. \cu, \gh, \qq\ and \aee. The three identified horseshoe librators \ql, \sa\ and \Larsen\ have short predictability period mostly due to the specific dynamic of horseshoe orbit. When an object is moving on a HS orbit, it has large amplitude of libration and experiences close multiple encounters with the planet (even if its eccentricity is small). This generates instability of the object's orbit and one can expect that such an asteroid is expelled from the co-orbital region after a few HS libration periods (see also discussion in Stacey and Connors (2008), section 4.1 therein). Indeed, according to our analysis these three objects execute a few HS loops (in the case of \ql\ and \Larsen\ the HS state is interrupted by QS behavior) before leaving this state, as well as the planet's co-orbital region. The comets \ar\ and \wc\ do not experiences a close encounters with Jupiter. On the other hand, these objects can close approach to Mars. Additionally, these comets have very short observational intervals. For these reasons the predictability period of these object is short. Asteroids \hr\ and \ug\ experience close encounters with Jupiter and they have relatively short predictability time interval. These objects also remain in the co-orbital region at most a few hundred of years. We conclude that the objects that experience close encounters with the Jupiter may be delivered into the co-orbital region from different parts of the Solar System, for example from either the Trojan area of external part of the System and they leave the co-orbital region relatively quickly (within a few hundred years). Kortenkamp and Joseph (2011) found that the Jupiter Trojans were able to migrate to QS orbits during the late stages of migration of the giant planets, however, it seems that current QS of this planet entered into this state later. The main reason is that the mechanism of transforming TP orbit into QS orbit was not efficient (about 0.2\% of the escaping Jupiter Trojan particles entered in QS after leaving the Trojan region within a few milion of years) and the QS of Jupiter can persist in this state at most $10^7$ yrs (Wiegert et al. 2000). Thus, the present QS entered into this state likely a few millions year ago or less. We have shown that \gh\ is currently a QS of Jupiter for about 200 years and it leaves this regime about 2030 transiting into TP orbit. In general, this object stays within the co-orbital region (all VOs remain in this area) and it jumps between QS and TP types of orbits. Similar dynamics demonstrates \gh, however, according to our calculations after 3900 it will be moving in TP-QS-TP orbit with large amplitude of the guiding center. This object has been observed for a short time interval of about 1.33 yr but it exhibits unexpectedly very long predictability period. The orbit of \gh\ is predictable to $\sim$4500. After that time the object is still locked in a 1:1 mean motion resonance with Jupiter to the end of integration. In this paper the detailed analysis of the four QS objects, i.e. two asteroids (\qq\ and \aee) and two comets (\ar\ and \wc) found by Kinoshita and Nakai (2007) have been presented. For both asteroids our results fully agree with the conclusions drawn by Kinoshita and Nakai (2007). This is especially worth emphazing because asteroid 2001 QQ199 was observed after the year 2007, thus for this object we used significantly richer observational material than Kinoshita and Nakai (2007). Since the NG effects may affect the dynamical behavior of comets \ars\ and \wcs\, we decided to integrate their orbits only 1000 years backward and forward. Within this period of integration all VOs of the first comet does not diverge significantly and all VOs shows QS behavior. In the case of \wcs\ we have shown that this object is a QS of Jupiter to the year 2500 at least. After that time the orbital elements of this comet starts to diverge and according to presented calculations the transitions between QS and TP orbit can take place. The main reasons why the predictability period of this object is very short seems to be that the comet has only about 0.5 yr interval of observations and the large eccentricity. Although \wcs\ does not pass close to Jupiter (all encounters are over 2.0 a.u.), it is able to experience very close encounters with Mars and the Earth due to the very eccentric orbit (e=0.68). We have confirmed the results of Karlsson (2004) that \hr\ will be temporarily captured by Jupiter on a QS orbit near 2050 and that the object \ug\ will be captured into co-orbital region of Jupiter around 2045, however, in contrary to his predictions, this asteroid is not co-orbital of Jupiter, presently. Its guiding center encloses the origin but this object will not a QS of Jupiter. Currently, these objects do not librate in a 1:1 mean motion resonance with Jupiter. This behavior causes that \hr\ and \ug\ repeatedly are coming close to Jupiter and therefore their predictability period is of order of a few hundred years. \newpage \Acknow{Part of numerical calculations was performed using the numerical orbital package developed by Professor Grzegorz Sitarski and the Solar System Dynamics \& Planetology Group at SRC PAS.} | 12 | 6 | 1206.0557 |
1206 | 1206.2078_arXiv.txt | A ring-shaped debris disk around the G2V star HD 202628 ($d$ = 24.4 pc) was imaged in scattered light at visible wavelengths using the coronagraphic mode of the Space Telescope Imaging Spectrograph on the {\it Hubble Space Telescope}. The ring is inclined by $\sim$64\deg\ from face-on, based on the apparent major/minor axis ratio, with the major axis aligned along PA = 130\deg. It has inner and outer radii ($>50$\% maximum surface brightness) of 139 AU and 193 AU in the northwest ansae and 161 AU and 223 AU in the southeast ($\Delta r/r \approx 0.4$). The maximum visible radial extent is $\sim254$ AU. With a mean surface brightnesses of V $\approx$ 24 mag arcsec$^{-2}$, this is the faintest debris disk observed to date in reflected light. The center of the ring appears offset from the star by $\sim$28 AU (deprojected). An ellipse fit to the inner edge has an eccentricity of 0.18 and $a$ = 158 AU. This offset, along with the relatively sharp inner edge of the ring, suggests the influence of a planetary-mass companion. There is a strong similarity with the debris ring around Fomalhaut, though HD 202628 is a more mature star with an estimated age of about 2 Gyr. We also provide surface brightness limits for nine other stars in our study with strong {\it Spitzer} excesses around which no debris disks were detected in scattered light (HD 377, HD 7590, HD 38858, HD 45184, HD 73350, HD 135599, HD 145229, HD 187897, and HD 201219). | Circumstellar debris disks, clouds of dust created by collisions of planetesimals such as comets and asteroids, provide evidence that stars are the central hosts of dynamic systems. Disk structures such as clearings, gaps, and asymmetrical dust distributions may indirectly reveal the presence of planets. The most prominent example is the Fomalhaut debris disk (Kalas et al.\ 2005) and its apparent planetary-mass companion, Fomalhaut b (Kalas et al.\ 2008; Chiang et al.\ 2009). The presence of dust around a star is most easily discerned by a measured infrared excess, though the disk is usually unresolved. While still relatively rare, resolved images of nearby debris disks spanning the millimeter to visible wavelength range have substantially increased in recent years, especially with the use of the {\it Hubble}, {\it Spitzer}, and {\it Herschel} space telescopes. The dust is often concentrated in a ring, as with Fomalhaut (Kalas et al.\ 2005), HD 4796 (Schneider et al.\ 2009), and HD 207129 (Krist et al.\ 2010). This makes detection easier because the surface brightness is greater than it would be in a more extended disk (e.g., $\beta$ Pictoris). In an effort to expand the number of resolved debris disks, which so far number around 20, we undertook an imaging survey using the {\it Hubble Space Telescope} ($HST$). For a list of candidate stars with {\it Spitzer}-measured infrared excesses and emission ratios of $L_{dust}/L_{star} > 10^{-4}$ we predicted the detectability of a ring-shaped disk around each assuming a width of $\Delta r/r = 0.2$, an albedo of 0.1, a canonical grain size distribution, and fractional scattered light brightness based on the $L_{dust}/L_{star}$. Ten nearby, solar-type targets with predicted inner radii outside of the inner working angle of the $HST$ Space Telescope Imaging Spectrograph (STIS) coronagraph were chosen. Our goal was to combine the dust distribution seen in the $HST$ images with {\it Spitzer} measurements to derive the actual grain properties (spatial distribution, albedo, scattering phase function, grain size distribution, etc.). Of the ten stars in our program, we have imaged a disk around only one, HD 202628 (HIP 105184, GJ 825.2, SAO 230622), a G2V star (V = 6.75) located at 24.4 pc. Koerner et al.\ (2010) measured a $\lambda$ = 70 $\mu$m excess ($17\times$ the photosphere) using {\it Spitzer}, deriving $L_{dust} / L_{star} = 1.4 \times 10^{-4}$. A central clearing is indicated by the lack of any 24 $\mu$m excess that would be emitted by warmer dust near the star. Using the assumptions discussed above, we predicted an inner disk radius of 80 AU (3\farcs 4). | HD 202628 is one of only three stars, including Fomalhaut and HR 4796, with eccentric, ring-shaped debris disks with sharp inner edges, all signs of planetary tidal interactions. It has the largest and most eccentric ring of the three. It is also the only solar-type star and is much older ($\sim2$ Gyr) than the other two. Based on the previous modeling of the Fomalhaut ring (Quillen 2006; Chiang et al. 2009), the perturbing planet in the HD 202628 system is probably close to, and inside of, the $\sim158$ AU inner edge. This shows that planets at very large orbital radii are present around not only more massive stars like Fomalhaut, which had presumably more massive accretion disks. This may be evidence that planets may form closer to the stars and get scattered to large distances by other, more massive planets. We hope to obtain deeper exposures of HD 202628 with $HST$ that would show this faint disk with better definition. These would allow more accurate characterizations of the inner edge and the outer extent of the disk. These will allow the use of modeling, as in Chiang et al. 2009, to more precisely constrain the location and mass of the perturbing planet. | 12 | 6 | 1206.2078 |
1206 | 1206.0817_arXiv.txt | {This paper investigate what is the main driver of the dust mass growth in the interstellar medium (ISM) by using a chemical evolution model of galaxy with metals (elements heavier than helium) in dust phase in addition to the total amount of metals. We consider asymptotic giant branch (AGB) stars, type II supernovae (SNe~II) and the dust mass growth in the ISM as the sources of dust, and SN shocks as the destruction mechanism of dust. Further, to describe the dust evolution precisely, our model takes into account the age and metallicity (the ratio of metal mass to ISM mass) dependence of the sources of dust. We particularly focused on the dust mass growth, and found that the dust mass growth in the ISM is regulated by the metallicity. To quantify this aspect, we introduce a ``critical metallicity'', which is a metallicity at which the contribution of stars (AGB stars and SNe~II) equals that of the dust mass growth in the ISM. If the star formation timescale is shorter, the value of the critical metallicity is higher, but the galactic age at which the metallicity reaches the critical metallicity is shorter. From observations, it was expected that the dust mass growth was the dominant source of dust in the Milky Way and dusty QSOs at high redshifts. By introducing the critical metallicity, it is clearly shown that the dust mass growth is the main source of dust in such galaxies with various star formation timescales and ages. The dust mass growth in the ISM is regulated by metallicity, and we stress that the critical metallicity works as an indicator to judge whether the grain growth in the ISM is dominant source of dust in a galaxy, especially because of the strong and nonlinear dependence on the metallicity.} | % Stellar light, in particular at shorter wavelengths, is absorbed by dust and re-emitted as the far-infrared thermal emission from the dust (e.g., Witt \& Gordon, 2000, and references therein). Therefore dust affects the spectral energy distributions of galaxies (e.g., Takagi et al., 1999; Granato et al., 2000; Noll et al., 2009; Popescu et al., 2011). The existence of dust in galaxies also affects the star formation activity. Dust grains increase the molecular formation rate by two orders of magnitude compared to the case without dust (e.g., Hollenbach \& McKee, 1979), and the interstellar medium (ISM) is cooled efficiently by molecules and dust. Consequently, star formation is activated drastically by dust. Hence, dust is one of the most important factors for the evolution of galaxies (e.g., Hirashita \& Ferrara, 2002; Yamasawa et al., 2011). \footnotetext{${}^*$the ratio of metal (elements heavier than helium) mass to ISM mass} The amount of dust in galaxies is one of the crucial factors to interpret the observational information of galaxies, since dust exists ubiquitously and the radiation from stars is always affected by dust attenuation. However, in spite of its importance, the evolution of dust amount has not been completely established yet. There are some key factors to understand dust evolution of galaxies. One of the keys is the ratio of metal (elements heavier than helium) mass to ISM mass, which is called ``metallicity''. Since dust grains consist of metals, it is natural to think that the evolution of dust is closely related to metallicity. In general, galaxies are believed to evolve from the state with a very low metallicity and very small amount of dust to higher amounts of metal and dust. Hence, it is a mandatory to model the formation and evolution of dust grains in galaxies along with the evolution of metallicity (e.g., Dwek \& Scalo, 1980; Hirashita, 1999a,b; Inoue, 2003; Yamasawa et al., 2011). Dust grains are formed by the condensation of metals. A significant part of the metals released by stellar mass loss during stellar evolution or supernovae (SNe) at the end of the life of stars condense into dust grains. Dust grains are not only supplied from stars but also destroyed by SNe blast waves (e.g., Jones et al., 1994; Jones, Tielens \& Hollenbach, 1996; Nozawa et al., 2003; Zhukovska, Gail \& Trieloff, 2008). In addition, we should consider the dust mass growth in the ISM by the accretion of atoms and molecules of refractory elements onto grains (e.g., Dwek, 1998; Liffman \& Clayton, 1989; Draine, 2009; Jones \& Nuth, 2011). What kind of dust formation processes are dominant at each stage of galaxy evolution? It is a very important question for understanding the evolution history of the ISM and star formation in galaxies. However, since dust evolution depends strongly on the age and metallicity of a galaxy, it is not easy to answer this question. Up to now, dust evolution has been studied with various models. For example, in young galaxies, SNe have been considered as the source of dust because they are the final stage of massive stars whose lifetime is short, and asymptotic giant branch (AGB) stars have been neglected because of their longer lifetime. However, Valiante et al. (2009) showed that the AGB stars also contribute to the dust production in young galaxies and cannot be neglected even on a short timescale of $\sim 500$~Myr. A more elaborate survey of the parameter space for the dust formation by SNe and AGB stars has been done by Gall et al. (2011a). They showed that the contribution of AGB stars exceeds that of SNe~II at several $100$~Myr if the ratio between metal and dust mass produced by SNe~II is less $\sim$ 0.01 and mass-heavy IMF with mass range $1$--$100\;\mbox{M}_{\odot}$. As for the dust mass growth, the ISM is considered to be the main source of dust in various galaxies. For example, the present dust amount observed in the Milky Way cannot be explained if the source of dust would have been only stars, suggesting that we must consider the dust mass growth in the ISM in evolved galaxies (e.g., Dwek, 1998; Liffman \& Clayton, 1989; Draine, 2009; Jones \& Nuth, 2011). Recently, dusty quasars (total dust mass $> 10^8\;{\rm M}_{\odot}$) have been discovered at high redshifts (e.g., Beelen et al., 2006; Wang et al., 2008), and theoretical studies on dust sources at high redshifts are currently carried out actively (e.g., Micha\l owski et al., 2010b; Gall et al., 2011a,b; Pipino et al., 2011; Valiante et al., 2011). They showed that it is hard to explain the total dust amount in these QSOs only with stellar contributions, and then they discussed the importance of the dust mass growth in the ISM. Next question is what controls the point where the dust mass growth in the ISM becomes efficient to the total dust mass in galaxies. Therefore, although each physical process has been already extensively discussed in preceding studies, there emerges a crucial question: what kind of dust production process is dominant at each stage of galaxy evolution? Particularly, when the dust dust mass growth becomes dominant as a source of dust mass? The central aim of this work is to address this question. In this paper, we investigate what is the main driver of the dust mass growth in the ISM. Since all sources of dust production are tightly related to each other on dust evolution, it is crucial to treat these processes in a unified framework to understand the evolution of dust in galaxies. Here, we adopt the model based on chemical evolution model in a same manner as Hirashita (1999b); Calura, Pipino \& Matteucci (2008); Inoue (2011). This is because their models consider main dust production/destruction processes that affect the dust evolution of galaxies, and it is easy to compare our results to the previous ones. {}From this work, we find that the dust mass growth in the ISM is regulated by metallicity. We refer to this metallicity as the critical metallicity, whose details are described in Sect.~3.2. Although the dust mass growth can occur at any time of a galaxy age, but we stress this point because there is a moment at which the dust mass growth overwhelms the contribution from other sources of dust. This paper is organized as follows. In Sect.~2, we describe the model developed for this work. In Sect.~3, we show and discuss the basic results obtained by our model. The main topic of this paper, critical metallicity, is introduced and extensively examined in Sect.~3.2. Section~4 is devoted to the conclusions. The solar metallicity is set to be ${\rm Z}_{\odot} = 0.02$ (Anders \& Grevesse, 1989) throughout this paper. | In this work, we constructed a galaxy evolution model taking into account the metallicity and age dependence on the various dust sources (AGB stars, SNe~II and growth in the ISM) to investigate what is the main driver of the grain growth which is expected to be the dominant source of dust in various galaxies with various star formation timescales. We have found that the timing that the dust mass growth in the ISM becomes effective is determined by metallicity. If metallicity in a galaxy exceeds a certain critical value, \textit{critical metallicity}, the dust mass growth becomes active and the dust mass rapidly increases until metals are depleted from the ISM. This critical metallicity is larger for a shorter star formation timescale. The dust mass growth is thought to be the dominant source of dust in evolved galaxies like the Milky Way and young but dusty and massive QSOs at high redshifts. The importance of the dust mass growth in such a diversity of galaxies can be explained clearly in terms of the critical metallicity; the dust mass growth in the ISM is regulated by metallicity, and we stress that the critical metallicity works as an indicator to judge whether the grain growth in the ISM is dominant source of dust in a galaxy, especially because of the strong and nonlinear dependence on the metallicity. | 12 | 6 | 1206.0817 |
1206 | 1206.5435_arXiv.txt | Results are presented for [CII] 158 \um line fluxes observed with the $Herschel$ PACS instrument in 112 sources with both starburst and AGN classifications, of which 102 sources have confident detections. Results are compared with mid-infrared spectra from the $Spitzer$ Infrared Spectrometer and with $L_{ir}$ from IRAS fluxes; AGN/starburst classifications are determined from equivalent width of the 6.2 \um PAH feature. It is found that the [CII] line flux correlates closely with the flux of the 11.3 \um PAH feature independent of AGN/starburst classification, log [f([CII] 158 \ums)/f(11.3 \um PAH)] = -0.22 $\pm$ 0.25. It is concluded that [CII] line flux measures the photodissociation region associated with starbursts in the same fashion as the PAH feature. A calibration of star formation rate for the starburst component in any source having [CII] is derived comparing [CII] luminosity L([CII]) to $L_{ir}$ with the result that log SFR = log L([CII)]) - 7.08 $\pm$ 0.3, for SFR in \mdot~ and L([CII]) in \ldot. The decreasing ratio of L([CII]) to $L_{ir}$ in more luminous sources (the ``[CII] deficit") is shown to be a consequence of the dominant contribution to $L_{ir}$ arising from a luminous AGN component because the sources with largest $L_{ir}$ and smallest L([CII])/$L_{ir}$ are AGN. | Understanding the initial formation of galaxies depends on discovering sources obscured by dust and tracing these sources to their earliest epoch in the universe. The extreme luminosity of dusty, local sources was originally revealed by the Ultraluminous Infrared Galaxies (ULIRGs, e.g. Soifer, Neugebauer and Houck 1987, Sanders and Mirabel 1996), whose luminosity arises from infrared emission by dust, and this dust often obscures the primary optical sources of luminosity. That such galaxies are important in the early universe was demonstrated by source modeling which indicated that the infrared dust emission from galaxies dominates the cosmic background luminosity \citep{cha01,lag04,lef05}. Surveys in the submillimeter were the first to discover individual, optically obscured, dusty sources at redshifts z $\ga$ 2 \citep{cha05}. A variety of observing programs using spectra from the $Spitzer$ Infrared Spectrometer (IRS; Houck et al. 2004) subsequently found luminous ULIRGS to redshifts z $\sim$ 3 \citep[e.g.][]{hou05,yan07,saj07,wee09b}. This $Spitzer$-discovered population of high redshift ULIRGs has large infrared to optical flux ratios [$f_{\nu}$(24 \ums) $>$ 1 mJy and $R$ $>$ 24] attributed to heavy extinction by dust and has been labeled ``dust obscured galaxies" (DOGS; Dey et al. 2008). Some DOGs are powered primarily by starbursts and some by active galactic nuclei (AGN), and the DOGS are similar to the population of submillimeter galaxies in overall spectral energy distributions (SEDs), redshifts, and luminosities \citep{pop08,men09,cop10,kov10}. To discover and understand dusty galaxies at even higher redshifts than the DOGs known so far, the atomic line emission of [CII] 158 \micron~is the single most important spectroscopic feature because it is the strongest far-infrared line \citep{sta91,luh03, bra08}. As a consequence, this line will provide the best opportunity for redshift determinations and source diagnostics using submillimeter and millimeter spectroscopic observations. Already, [CII] has been detected at redshift exceeding 7 \citep{mai05,ven11} and shown to be strong in starbursts with 1 $<$ z $<$ 2.5 \citep{hai10,sta10,ivi10}. Our primary motives for the present paper are to present [CII] results for a large sample of dusty sources and to compare with mid-infrared classification indicators for starbursts and AGN. This comparison leads to a calibration between star formation rate (SFR) and [CII] luminosity. We emphasize the diagnostics used for DOGS at z $\sim$ 2, because the large populations of submillimeter and mid-infrared DOGS now known at this epoch provide a crucial reference for scaling to higher redshifts. The epoch 2 $\la$ z $\la$ 3 is also important because this is the observed epoch at which starburst and AGN activity seems to peak \citep[e.g.][]{mad98,red09,fan04,cro04,bro06}. The [CII] line should be primarily a diagnostic of star formation, being associated with the photodissociation region (PDR) surrounding starbursts \citep{tie85,hel01,mal01,mei07}, and the line appears to be weaker in the most luminous sources (``the [CII] deficit", Luhman et al. 2003). It is crucial to understand the origin of this line and the extent to which its luminosity is a measure of SFR. Does [CII] scale with other star formation indicators? Is the [CII] deficit a consequence of AGN dominance rather than star formation in luminous sources? Determining such answers is the objective of new observations we have undertaken with the $Herschel$ Space Observatory \citep{pil10}. | 12 | 6 | 1206.5435 |
|
1206 | 1206.2215_arXiv.txt | We present the orbital-phase resolved analysis of an archival FO Aqr observation obtained using the X-ray Multi-Mirror Mission (\emph{XMM-Newton}), European Photon Imaging Camera (pn instrument). We investigate the variation of the spin pulse amplitudes over the orbital period in order to account for the effects of orbital motion on spin modulation. The semi-amplitude variations are in phase with the orbital modulation, changing from (38.0$\pm$1.8)\% at the orbital maximum to (13.3$\pm$3.7)\% at the orbital minimum. The spectral parameters also show changes over the orbital period. One of the absorption components increase by a factor of 5 between the orbital minimum and maximum. We interpret that this absorption arises from the bulge where accretion stream from the secondary impacts the disk. The spectrum extracted from the orbital minima and maxima can be fitted with a warm absorber model yielding values $N_{\rm{H}}$ = 2.09$^{+0.98}_{-1.09}$$\times$ $10^{22}$ and 0.56$^{+0.26}_{-0.15}$$\times$ $10^{22}$ cm$^{-2}$; and log($\xi$) = 0.23$^{+0.37}_{-0.26}$ and $<$0.30 erg cm s$^{-1}$ respectively, indicating the existence of ionized absorption from the bulge at the impact zone which is spread out on the disk. The absorption due to accretion curtain and/or column which causes the spin modulation can be distinguished from the disk absorption via spectral modeling. | FO Aqr is a compact binary belonging to the sub-category intermediate polars (IPs) of the cataclysmic variables (CVs) which are composed of a white dwarf accreting material from a Roche lobe filling main sequence star. IPs have white dwarfs with a magnetic field of mostly 1-10 MG (although there are a couple of systems with magnetic field strengths up to 30 MG; e.g. Katajainen et al. 2010, Piirola et al. 2008). The accretion occurs through a truncated disk and via accretion curtains to the magnetic poles of the white dwarf (see Patterson 1994; Warner 2003). FO Aqr is a well studied IP, with an orbital period of 4.85 hr (Osborne \& Mukai 1989) and white dwarf spin period of 20.9 min (Patterson et al. 1998). It has been observed with almost every X-ray mission; \emph{EXOSAT}, \emph{Ginga}, \emph{ASCA}, \emph{RXTE}, \emph{XMM-Newton}, \emph{INTEGRAL}, \emph{SWIFT} and \emph{Suzaku} (Cook, Watson \& McHardy 1984, Norton et al. 1992; Mukai, Ishida \& Osborne 1994; Evans et al. 2004, Parker, Norton \& Mukai 2005, Landi et al. 2009, Yuasa et al. 2010, respectively). Our understanding of the accretion scenario in FO Aqr has changed over the years: Norton et al. (1992) suggested that the system shows diskless accretion and that the accretion flow changes poles. However, Hellier (1993) argued that there is an accretion disk in the system and the accretion takes place both via a disk and stream overflowing the disk. Mukai, Ishida \& Osborne (1994) confirmed this hybrid accretion mode, but proposed that accretion from a partial disk is dominant. Later, it was suggested that the accretion mode alternates from a hybrid of disk-fed and stream fed accretion to disk-fed accretion over the years (i.e. hybrid in 1988, disk-fed in 1990 (Norton, Beardmore \& Taylor 1996) hybrid in 1993 and 1998 (Beardmore et al. 1998), disk-fed in 2001 (Evans et al. 2004)). The changes in accretion modes are most likely due to changes in the mass accretion rate of the system (de Martino et al. 1999). The source shows orbital modulations, which are deeper at the lower energy regime, indicating that it could be due to absorption from structures on the disk and the accretion stream (Hellier et al. 1993; Evans et al. 2004; Parker, Norton \& Mukai 2005). The spin pulse shape of the system is complicated, with a quasi-sinusoiadal component and a notch after a "dip" caused by the accretion curtains (Evans et al. 2004). The spin pulse profile may vary between nearly sinusodial to saw-tooth shapes in different observations (Beardmore et al. 1998). The X-ray spectrum of the source can be represented with a multiple plasma emission component, complex absorption and Gaussian lines (eg. Mukai, Ishida \& Osborne 1994; Evans et al. 2004; Yuasa et al. 2010). A soft blackbody emission was not detected by Evans \& Hellier (2007) with \emph{XMM-Newton} data or by Yuasa et al. (2010) with \emph{Suzaku} data. However, using the results of a combined joint analysis of \emph{INTEGRAL/IBIS} and \emph{SWIFT/XRT} data, Landi et al. (2009) claimed the detection of a blackbody component with a temperature of 61 eV which they attributed to the irradiation of the white dwarf atmosphere. In this work, the spectral and temporal properties of FO Aqr over the orbital phase will be investigated. In Section 2 observation and data preparation will be introduced. In Section 3 the spin pulse profile behavior over the orbital phase will be outlined. In Section 4 the phase average spectrum and the variation of spectral parameters over the orbital phase will be investigated. In Section 5 the results of the analyses will be discussed. | We have presented X-ray orbital-phase resolved analysis of the intermediate polar FO Aqr. This work improves upon the previous works by calculating the spectral parameters over the orbital phase. The distinction between the absorbing components are clarified and the values are explicitly calculated. The absorption originating from the polar regions of the white dwarf can be resolved from the absorption by structures on the accretion disk. The shape of the spin pulse profile is unaffected by the orbital motion, however the semi-amplitude of the profile change over the orbital phase. The X-ray orbital variation over the orbit in the system arises from absorption by the bulge material on the disk spread well over the disk. The absorption column is greatly enhanced during the X-ray orbital dip. Moreover, we have modeled the absorption from the orbital dip with a warm absorber model for the first time for this source and also for CVs, confirming the ionized nature of the material on the disk causing the absorption. We derived a range of ionization parameter log($\xi$$_{CV}$)$\sim$ -0.8--1 for plausible warm absorbers on CV disks. We calculate that the orbital modulations in FO Aqr (and plausibly some other CVs) are similar to those seen in LMXB dippers as variations (density/temperature) of a warm absorber at the accretion impact region and spreading on the disk. | 12 | 6 | 1206.2215 |
1206 | 1206.0483_arXiv.txt | We present previously unpublished July 2005 $H$-band coronagraphic data of the young, planet-hosting star HR 8799 from the newly-released Keck/NIRC2 archive. Despite poor observing conditions, we detect three of the planets (HR 8799 bcd), two of them (HR 8799 bc) without advanced image processing. Comparing these data with previously published 1998-2011 astrometry and that from re-reduced October 2010 Keck data constrains the orbits of the planets. Analyzing the planets' astrometry separately, HR 8799 d's orbit is likely inclined at least 25$^\circ$ from face-on and the others may be on in inclined orbits. For semimajor axis ratios consistent with a 4:2:1 mean-motion resonance, our analysis yields precise values for HR 8799 bcd's orbital parameters and strictly constrains the planets' eccentricities to be less than 0.18--0.3. However, we find no acceptable orbital solutions with this resonance that place the planets in face-on orbits; HR 8799 d shows the largest deviation from such orbits. Moreover, few orbits make HR 8799 d coplanar with b and c, whereas dynamical stability analyses used to constrain the planets' masses typically assume coplanar and/or face-on orbits. This paper illustrates the significant science gain enabled with the release of the NIRC2 archive. | The nearby, young A-type star HR 8799 \citep[$d$ = 39.4 $pc$, $\approx$ 30 Myr;][]{Zuckerman2011} harbors the first independently confirmed, directly imaged exoplanetary system and the only imaged multi-planet system \citep{Marois2008}. After the discovery of HR 8799 bcd ($r_\mathrm{proj}$ $\approx$ 24, 38, and 68 AU) reported in November 2008 \citep{Marois2008}, other studies identified at least one of these planets in archival data taken prior to 2008 \citep{Lafreniere2009,Fukagawa2009,Metchev2009,Soummer2011}. HR 8799 planet astrometry derived from both pre and post-discovery images can help constrain the system's dynamical stability and, in turn, the planets' physical properties. At least two of the HR 8799 planets are likely locked in a mean motion resonance, otherwise the system would quickly become dynamically unstable \citep{Fabrycky2010}. The recently discovered fourth companion at $\sim$ 15 AU, HR 8799 e, generally makes dynamical stability less likely \citep{Marois2011,Currie2011a}, favoring lower masses of $M_\mathrm{b,cde}$ $<$ 7, 10 $M_\mathrm{J}$, an important constraint given the uncertainties in deriving masses from planet cooling and atmosphere models \citep[][]{Spiegel2012,Madhusudhan2011}. Studies focused on fitting the planets' orbits and/or testing dynamical stability typically assume that the planets are a) in resonance (4:2:1 for HR 8799 bcd or 2:1 for HR 8799 cd), b) in circular, face-on orbits, c) and/or in coplanar orbits \citep[e.g.][see also Fabrycky and Murray-Clay 2010]{Marois2011,Currie2011a}. However, \citet{Soummer2011} show that circular, face-on, and coplanar orbits are inconsistent with 1998 HST astrometry, identifying a best-fit orbit for HR 8799 d of $i$ = 28$^\circ$ and $e$ = 0.115. Generally, more eccentric orbits destabilize the system. The system stability depends on the (mutual) inclinations of the planets \citep{Fabrycky2010,Sudol2012}. Thus, the HR 8799 planets' true mass limits derived from dynamical stability arguments may slightly differ from those previously reported. Well-sampled HR 8799 d astrometry could help clarify whether HR 8799 d's orbit must be inclined, eccentric, and/or coplanar with the other planets. However, until now there is a $\sim$ 9-year gap between the 1998 HST detection and the next one (2007; Metchev et al. 2009). New astrometry for HR 8799 bce in between 1998 and 2007 could also help constrain those planets' orbits. By better determining the HR 8799 planets' orbital properties, we can more conclusively investigate system dynamical stability and thus better clarify the range of allowable planet masses. In this Letter, we report the detection of HR 8799 bcd from unpublished, now-public Keck/NIRC2 data taken in 2005 supplemented with a re-reduction of published October 2010 data from \citet{Marois2011}. We use these data to better constrain the orbital properties of HR 8799 bcd. | From analyzing HR 8799 bcd astrometry from our new ``pre-discovery" image and other data, we provide new constraints on the planets' orbital properties. Treating the three planets separately, we narrowly constrain three major orbital parameters ($a$/$i$/$e$) for HR 8799 c and d. None of the planets are likely to be orbiting face-on and the inclinations for acceptably-fitting orbits are systematically higher for HR 8799 d than for HR 8799 b and c. If HR 8799 bcd have semimajor axes consistent with a 4:2:1 resonance, our analysis strongly constrains the major orbital properties for all three planets. The three planets (especially c and d) then even more obviously have inclined orbits. Most acceptable solutions for HR 8799 d place the planet on an orbit inclined by more than 7$^\circ$ (21$^\circ$) relative to HR 8799 b(c)'s orbit: few orbital solutions consistent with the astrometry also place them on coplanar orbits. Adopting a less restrictive definition for ``acceptably-fitting" orbits does not undo any of these trends, although there are more orbit combinations making the planets coplanar. Adopting the median parameter value or MLO instead of the more conservative 68\% confidence interval likewise does not change these results. These results provide valuable input for constraining the mass of the HR 8799 planetary system. Longer-term astrometric monitoring of HR 8799 \citep[i.e.][]{Konopacky2011} will better clarify the planets' orbital properties. Limits on the planets' dynamical masses will provide crucial input for planet cooling models and even more firmly establish HR 8799 as a benchmark system to understand the properties of young, self-luminous planets. Finally, this work and other recent studies of HR 8799 \citep{Soummer2011,Lafreniere2009,Fukagawa2009} clearly demonstrate the value of publicly archiving data on advanced telescopes. In our case, detecting at least two HR 8799 planets (HR 8799 bc) was rather straightforward and did not require advanced image processing techniques developed well after the data were taken. As data for Keck and many other 8--10 m class telescopes are now archived, they provide an indispensible resource with which to confirm and characterize directly imaged planets like HR 8799's and other substellar companions. | 12 | 6 | 1206.0483 |
1206 | 1206.2165_arXiv.txt | {With now more than 20 exoplanets discovered by CoRoT, it has often been considered strange that so many of them are orbiting F-stars, and so few of them K or M-stars. Up to now, studies of the relation between the frequency of extrasolar planets and the spectral types, or masses of their host stars has been the realm of radial velocity surveys. Although transit search programs are mostly sensitive to short-period planets, they are ideal for verifying these results. This is because transit search programs have different selection biases than radial velocity surveys. To determine the frequency of planets as a function of stellar mass, we also have to characterize the sample of stars that was observed.} {We study the stellar content of the CoRoT-fields IRa01, LRa01 (=LRa06), and LRa02 by determining the spectral types of 11\,466 stars. Nine planet-host stars have already been identified in these fields. Determing the spectral types of thousands of stars of which CoRoT obtained high-precision light-curves also makes possible a wide variety of other research projects.} {We used spectra obtained with the multi-object spectrograph AAOmega and derived the spectral types by using template spectra with well-known parameters.} {We find that $\mathrm{34.8\pm0.7\%}$ of the stars observed by CoRoT in these fields are F-dwarfs, $\mathrm{15.1\pm0.5\%}$ G-dwarfs, and $5.0\pm0.3\%$ K-dwarfs. We conclude that the apparent lack of exoplanets of K- and M-stars is explained by the relatively small number of these stars in the observed sample. We also show that the apparently large number of planets orbiting F-stars is similarly explained by the large number of such stars in these fields. Given the number of F-stars, we would have expected to find even more F-stars with planets . Our study also shows that the difference between the sample of stars that CoRoT observes and a sample of randomly selected stars is relatively small, and that the yield of CoRoT specifically is the detection one hot Jupiter amongst $\mathrm{2100\pm700}$ stars. } {We conclude that transit search programs can be used to study the relation between the frequency of planets and the mass of the host stars, and that the results obtained so far generally agree with those of radial velocity programs. } | Analyzing the statistical properties of extrasolar planets provides important clues on planet-formation. Particularly important are the relations between the properties of host stars and their planets. Up to now, most of these studies have been carried out using radial velocity surveys. Radial velocity surveys have the clear advantage that they allow one to detect planets with orbital periods ranging from less than a day to many years, but they have the disadvantage that they are biased against rapidly rotating and active stars. This is a problem if one aims to determine the relation between the mass of the host stars and the frequency of planets, because most of the main-sequence stars with masses larger than the Sun rotate fast (e.g. Reiners \& Schmitt \cite{reiners03}). One possible solution is to observe giant stars. However, giant stars do not have close-in planets and it is therefore difficult to compare the results obtained for giant stars with those for main-sequence stars. Nevertheless, by combining data of main-sequence and giant stars, Johnson et al. (\cite{johnson10}) concluded that the frequency of planets increases proportionally to the mass of the host stars (see also Johnson et al. \cite{johnson07}; Udry \& Santos \cite{udry07}). It would consequently be important to test these results with an independent method. Although transit search programs detect mostly short-period planets, they offer this possibility. The CoRoT-survey (COnvection, ROtation, and planetary Transits) is particularly suitable for this purpose. The discovery of a planet transiting an F6V star with a $v\,\sin\,i$ of $40\pm5$ $\mathrm{km\,s^{-1}}$ (Gandolfi et al. \cite{gandolfi10}), and a planet orbiting a highly active star (Alonso et al. \cite{alonso08}) shows that this survey is not biased against rapidly rotating and active stars. Furthermore, the detection of a planet with an orbital period of 95 days (Deeg et al. \cite{deeg10}), and a planet with radius of $1.58\pm0.10\, R_{\rm Earth}$ (L\'eger et al. \cite{leger09}; Bruntt \cite{bruntt10}) shows that CoRoT is able to detect planets in relatively long orbit as well as planets of small radii. The statistical analysis of the CoRoT survey thus will give us a complementary view to the results obtained in radial velocity surveys, because it is less biased against rapidly rotating and active stars. As an additional advantage, CoRoT observes fields in different regions of the sky, allowing us to find out whether the population of planets depends on the region in the sky that is being observed. At first glance there seems to be a huge difference between results obtained in transit and radial velocity surveys. Out of the first 20 planet-hosting stars discovered by CoRoT, 35\% are F-stars, 55\% G-stars, and 20\% are K-stars, with not a single one being an M-star. If we consider the number of stars harboring planets with a semi-major axis $\leq0.1$ AU discovered by radial velocity surveys, 7\% of them are F-stars, 39\% G-stars, 39\% K-stars, and 13\% M-stars. It has therefore often been considered strange that CoRoT finds so many planets orbiting F-stars, and so few orbiting K and M-stars. However, before we can conclude that there is a difference between the two types of surveys, we have to know how many F, G, K and M-stars the samples contains. To compare the results obtained in radial velocity surveys with those obtained by CoRoT, we have to know how many F, G, and K-type stars CoRoT has observed. This is the aim of this work. By spectroscopically characterizing the sample of stars that CoRoT has observed, we answer the question whether relatively few planets are found around K- and M-stars because of the lack of such objects in the sample, or whether it reflects a real lack of close-in planets around these types of stars. In the same way, we also answer the question why CoRoT finds so many planets of F-stars. | Of these nine planet-host stars discovered in the three CoRoT-fields studied in this work, four are F-type stars (\object{CoRoT-4}, Moutou et al. \cite{moutou08}; \object{CoRoT-5}, Rauer et al. \cite{rauer09}; \object{CoRoT-14}, Tingley et al. \cite{tingley11}; \object{CoRoT-21} P\"atzold et al. \cite{paetzold11}), four are G-type stars (\object{CoRoT-1}, Barge et al. \cite{barge08}; \object{CoRoT-7}, L\'eger et al. \cite{leger09}; \object{CoRoT-12}, Gillon et al. \cite{gillon10}; \object{CoRoT-13}, Cabrera et al. \cite{cabrera08}), and one is a K-star (\object{CoRoT-24}; Alonso et al. \cite{alonso12}). With only nine planet-host stars, the sample is too small to answer the question whether or not the number of planets increases with the mass of the host star. However, we can still discuss whether the high percentage of planets orbiting F-stars and the small fraction orbiting K-stars can be explained with the observed sample of stars, or whether it can already be taken as evidence for an increase of the planet frequency with the mass of the host star. Given that only $5.0\pm 0.3\%$ of the stars in the sample are K- and M-dwarfs, the small number of planets found in K- and M-stars by CoRoT is due to the lack of K- and M-stars in the sample. The situation is more complicated for the F-stars. Since four of the planet-hosting stars are G-stars, and there are 2.3 times as many F- than G-stars, there should have been 9 to 14 planet-hosting F-stars, depending on whether we assume that the frequency of planets increases with the mass of the star, or not. In contrast to this, only four planet-hosting stars of this type were found. Of course, this is still low-number statistics, but we can conclude that there is no excess of close-in F-star planets. Quite contrary to this, we would have expected to find even more planets of F-stars. If we add up the results obtained for all galactic anti-center fields, CoRoT has found four F-stars and eight G-stars that are hosting planets. If there were also twice as many F- as G-stars in these fields, we would have expected to find four times as many F-stars hosting planets than we have found. However, we have to carry out a similar study of all the other fields before we can draw any firm conclusion. In any case, our study shows that the statistical analysis of transit search programs can lead to very interesting results. Once the CoRoT survey is completed, it will be possible to answer the question whether or not there is a lack of close-in planets of F-stars. Up to now, we have only analyzed stars located in the so-called "galactic anti-center region" of CoRoT. The next step would be to carry out a similar survey for the fields in the so called "galactic center eye" of CoRoT. Apart from improving the statistics, there is another good reason: Because of the metallicity gradient in our galaxy and the correlation between planet frequency and metallicity, it is expected that stars orbiting closer to the galactic center should have a higher frequency of planets than stars orbiting at larger distances. If we take the gradient of $-0.06\pm0.01\,\mathrm{kpc^{-1}}$ from Friel et al. (\cite{friel02}) and the relation between metallicity and planet frequency from Santos et al. (\cite{santos04}), we expect a difference of a factor of two for the planet frequency of stars with galactocentric distances that differ by 2 kpc. Using a more sophisticated model, Reid (\cite{reid08}) estimated that the difference of the planet-frequency is a factor of 1.8 for stars with galactocentric distances between 7 and 9 kpc. Since CoRoT observes stars in this range CoRoT-data might show this effect, if all fields are analyzed. Interestingly, the galactocentric distance from 7 to 9 kpc corresponds also to the galactic habitable zone (Lineweaver \cite{lineweaver04}). By extending the survey to the "galactic center" fields we will thus not only improve the statics but we might even be able to find out whether the properties of the planets depend on the galactocentric distance. | 12 | 6 | 1206.2165 |
1206 | 1206.2679_arXiv.txt | We present the first systematic 1.4 GHz Very Large Array radio continuum survey of fossil galaxy group candidates. These are virialized systems believed to have assembled over a gigayear in the past through the merging of galaxy group members into a single, isolated, massive elliptical galaxy and featuring an extended hot X-ray halo. We use new photometric and spectroscopic data from SDSS Data Release 7 to determine that three of the candidates are clearly not fossil groups. Of the remaining 30 candidates, 67\% contain a radio-loud (L$_{1.4GHz} > 10^{23}$ W Hz$^{-1}$) active galactic nucleus (AGN) at the center of their dominant elliptical galaxy. We find a weak correlation between the radio luminosity of the AGN and the X-ray luminosity of the halo suggesting that the AGN contributes to energy deposition into the intragroup medium. We only find a correlation between the radio and optical luminosity of the central elliptical galaxy when we include X-ray selected, elliptically dominated non-fossil groups, indicating a weak relationship between AGN strength and the mass assembly history of the groups. The dominant elliptical galaxy of fossil groups is on average roughly an order of magnitude more luminous than normal group elliptical galaxies in optical, X-ray, and radio luminosities and our findings are consistent with previous results that the radio-loud fraction in elliptical galaxies is linked to the stellar mass of a population. The current level of activity in fossil groups suggests that AGN fueling continues long after the last major merger. We discuss several possibilities for fueling the AGN at the present epoch. | A ``fossil galaxy group'' is believed to be the merger remnant of a galaxy group whose most massive members have succumbed to dynamical friction. The low velocity dispersion of the group environment results in a high interaction rate between galaxies, such that the most massive members eventually coalesce to form a single giant elliptical. What remains is the faint dwarf galaxy population, and an extended, hot gaseous halo that fills the dark matter potential of the former group. As the dynamical end point of groups, fossil groups appear to be undisturbed, virialized systems. They are identified observationally by two criteria: fossil groups are dominated by a single elliptical galaxy, typically brighter than L$_*$, where the next brightest companion is fainter by at least two magnitudes in R-band; and an extended, hot X-ray halo, typically brighter than L$_X=10^{42}$ ergs s$^{-1}$ \citep{Jones03}. The first fossil group system, RX J13040.6+4018, was identified by comparing ROSAT extended deep survey images with Palomar sky survey plates \citep{Ponman94}. In the last decade, several authors have contributed to a growing list of known fossil groups (e.g. 3, \citealt{Vikhlinin99}; 5, \citealt{Jones03}; 3, \citealt{2004AdSpR..34.2525Y}; 1, \citealt{2004MNRAS.349.1240K}; 1, \citealt{Sun04}; 1, \citealt{2005ApJ...624..124U}). Recently, \cite{Santos07} used the Sloan Digital Sky Survey (SDSS) Data Release 5 (DR5) and the ROSAT All-Sky Survey (RASS) to catalog 34 fossil group candidates allowing for the beginning of statistical investigations into the nature of such objects. Since the undertaking of our radio survey, \cite{Eigenthaler09} reported 34 new fossil groups, and \cite{LaBarbera09} identified an additional 25 fossil group candidates in SDSS DR4 and RASS. The formation of fossil groups is inferred from the comparison of the optical luminosity of the central elliptical galaxy and the X-ray luminosity of the halo with observations of systems in a range of similar or intermediate evolutionary states. The optical and X-ray values of fossil groups are comparable to compact galaxy groups---systems believed to be in the most advanced stages of merging \citep{MendesdeOliveira07}. Further, they are more numerous than compact groups suggesting that the merging occurs on short timescales compared to the lifetime of the group (e.g. \citealt{Mamon87}). The end result is the largest galaxies in compact groups succumb to dynamical friction to coalesce into a single, massive elliptical galaxy \citep{MendesdeOliveira06,MendesdeOliveira07}. Cosmological N-body simulations suggest that fossil groups formed early in the lifetime of the Universe, assembling half their mass by redshift of $z\ge1$ \citep{D'onghia05}, or by $z\ge0.8$ \citep{vonBendaBeckmann08}. Observationally, the luminosity function, the large magnitude gap between the first and second ranked galaxies, and the large virialized X-ray halo are believed to be indicative of this early formation \citep{Ponman94,Jones03,MendesdeOliveira06}. However, simulations also show that fossil groups may only be a phase of hierarchical structure evolution which ends with the fresh infall of massive galaxies from the surrounding environment \citep{vonBendaBeckmann08,Dariush10}. It has been noted that the dominant elliptical galaxy in fossil groups frequently resembles that of the brightest cluster galaxies (e.g. \citealt{Vikhlinin99,MendesdeOliveira06}), and because fossil groups have a similar space density to galaxy clusters ($\sim$$10^{-6}$ Mpc$^{-3}$; \citealt{Jones03}; \citealt{Santos07}), it has been suggested that may be the early seeds around which present-day clusters form \citep{Jones03}. Fossil groups are most well studied at optical and X-ray wavelengths, and it has been demonstrated that a strong relationship exists between the X-ray and R-band luminosities of the hot halo and the central dominant elliptical galaxy, respectively. This relationship indicates that the conditions of the group environment are linked to the mass assembly history of the fossil group \citep{Jones03}. However, it is unclear what mechanism is responsible for elevating the X-ray temperature in fossil groups. Clusters and elliptically dominated non-fossil groups follow the well known X-ray luminosity-temperature relationship (e.g.~\citealt{White97,Jones00}) and the temperatures of these systems imply the presence of non-gravitational heating in the intracluster and intragroup medium (ICM/IGM). Radio active galactic nuclei (AGN) commonly occur in these systems and their detailed morphology frequently corresponds to cavities in the X-ray halo, thus they are invoked to account for excess energy in the ICM/IGM (e.g. \citealt{Birzan04,Wise07}). Fossil group X-ray halos for which the temperature has been measured appear to follow a similar $L_X-T_X$ relation although they are more luminous relative to non-fossils implying that if the X-ray luminosity is boosted, so is the temperature \citep{Khosroshahi07}. Among non-fossil groups, evidence shows that radio-loud groups are hotter in X-rays than radio-quiet groups, indicating that on-going nuclear activity contributes to heating the IGM \citep{Croston05}. Simulations show that radio jets from AGN may deposit significant amounts of energy in the IGM (e.g. \citealt{Nath02,Heinz06}), and both simulations and observations of clusters support the theory that radio AGN can provide a sufficiently distributed heating mechanism to account for the excess energy (\citealt{Bruggen02,Fabian03}). However, simulations by \citet{Rowley04} suggest that major mergers can also elevate the X-ray temperature and luminosity in clusters of a few $\times10^{14}$ M$_{\odot}$ and that this may last for up to a Gyr after the stellar merger signatures have faded, decreasing the necessity of AGN to account for excess energy in the IGM. A radio investigation is therefore important for understanding the role of radio AGN in heating the IGM in fossil groups. A multi-wavelength investigation will allow us to link the properties of the radio AGN (or lack thereof) to the group environment as traced by the X-ray halo, and mass assembly of the central galaxy as indicated by the $r$-band stellar luminosity. Further, if radio AGN are common in fossil groups, what physical mechanism is responsible for maintaining or triggering them long after the majority of mass assembly has taken place in these systems? There is some evidence that optical AGN activity is the result of major mergers at late times \citep{Alonso07}, however observations and simulations suggest the last major merger in fossil groups occurred well over 1 Gyr in the past \citep{D'onghia05,vonBendaBeckmann08,Jones00,2006MNRAS.372L..68K}. Additionally, radio AGN activity, or indeed any AGN activity, appears to be well removed from merger activity, particularly since a redshift of $z=1$ \citep{Schawinski10,Cisternas11}. The thermodynamic properties of the X-ray halo gas, undisturbed by recent mergers, exhibited by the smooth, symmetric distribution of the X-ray contours, make the fossil group an ideal observational laboratory in which to test theories of AGN feedback for regulating the temperature of group and cluster sized halos \citep{Jetha09}. In this paper, we present the first systematic radio study of fossil group candidates taken from the \citet{Santos07} ROSAT/SDSS DR5 catalog, with a modestly deep 1.4 GHz Very Large Array (VLA) survey, and use our results to conduct a multi-wavelength investigation of fossil groups. We measure the radio flux as a tracer of AGN activity and our results reveal the degree to which radio AGN are \textit{currently} contributing to the energy deposition of the IGM. We calculate the fraction of radio-loud AGN and compare the population of radio sources with those in elliptically dominated groups and clusters which represent the continuum of hierarchical structures. The presence of radio AGN indicate that physical mechanisms such as cooling flows, minor mergers, or late-time accretion, are important for sustaining an AGN late in the evolution of fossil group systems, long after the majority of mass assembly has taken place. Thus, our radio survey is strongly motivated by outstanding questions in the evolution of galaxies, gas accretion, and the hierarchical formation of large scale structure. The rest of the paper is organized as follows: in Section \ref{obs} we describe our VLA 1.4 GHz radio observations and the data reduction, and describe the acquisition of optical data from SDSS DR7. In Section \ref{results} we compare published X-ray data based on the \citet{Santos07} sample, with SDSS DR7 \emph{r}-band photometry and with our 1.4 GHz radio results. In Section \ref{discussion} we discuss the possibilities for fueling a radio source long after the last major merger with the central, dominant elliptical galaxy. Finally, we conclude in Section\ref{conclusions}. | \label{conclusions} We present the first systematic survey of fossil galaxy group candidates at radio wavelengths using one of the largest collection of uniformly identified sources in the literature \citep{Santos07}. With updated photometric and spectroscopic optical data from SDSS DR7, we find 30 of these 33 candidates are fossil groups and that 67\% of them are detected at 1.4 GHz down to an average \emph{rms} noise of $\sim$80 $\mu$Jy. All but one fossil group detection are radio-loud sources. This is a large fraction considering fossil groups were believed \emph{a priori} to be old, quiescently evolving galaxy systems. However, it is consistent with previous studies that correlate a rise in the fraction of radio-loud AGN with a rise in the stellar mass of the host galaxy \citep{Best05}. We find that MKW/AWM poor clusters \citep{Morgan75,Albert77} may be an excellent sample, spanning a range of $\Delta$mag between the first- and second-ranked galaxies, against which to compare fossil groups to further understand the history of their mass assembly. The high AGN frequency in fossil groups is most likely not directly related to their unique formation history given the prevalence of radio sources among X-ray selected samples of elliptical galaxies and non-fossil groups \citep{Croston05,Dunn10} and the short radiative lifetime of synchrotron emission at 1.4 GHz compared to the simulated formation time scale of fossil groups \citep{D'onghia05,vonBendaBeckmann08}. We find a weak trend between the radio and X-ray luminosities of fossil groups, but no trend between the radio and optical properties of the central elliptical galaxy. However, when we include the brightest group galaxy from a sample of X-ray selected elliptically dominated groups, we find both radio--X-ray and radio--optical correlations such that fossil groups are more luminous at all three wavelengths than non-fossil groups. Thus, if optical luminosity represents the mass assembly history, and the X-ray halo represents the group-sized dark matter halo in which the galaxy resides, then when an AGN is ``on'', its radio strength is related to the environment in which it resides. The dominant elliptical galaxy in a fossil group resides at the center of group gravitational potential revealed by the presence of an extended X-ray halo. In such an environment, the nuclear activity may tied to the thermodynamic state of the hot IGM gas and a cooling flow may both fuel and be regulated by the AGN. However, there is not yet direct evidence of a cooling flow in these systems. Deep follow-up imaging of the central galaxies at optical and radio wavelengths is required to exclude the possibility of recent mergers or the existence of nearby cold gas reservoirs that can fuel the central engine. Finally, fossil groups by definition harbor a population of nearby dwarf galaxies; it is unknown to what degree minor mergers contribute to feeding the black hole and the signatures may be difficult to detect. It is clear that fossil groups are not only an interesting study of mass assembly in the early Universe, but also as a source of accretion and feedback with the IGM in the present day: their isolation provides a relatively simplified environment in which to study these processes. | 12 | 6 | 1206.2679 |
1206 | 1206.7026_arXiv.txt | We present a study of Broad Absorption Line (BAL) quasar outflows that show \siv~$\lambda$1063 and \siv*~$\lambda$1073 troughs. The fractional abundance of \siv\ and \civ\ peak at similar value of the ionization parameter, implying that they arise from the same physical component of the outflow. Detection of the \siv* troughs will allow us to determine the distance to this gas with higher resolution and higher signal-to-noise spectra, therefore providing the distance and energetics of the ubiquitous \civ\ BAL outflows. In our bright sample of 156 SDSS quasars 14\% show \civ\ and 1.9\% \siv\ troughs, which is consistent with a fainter magnitude sample with twice as many objects. One object in the fainter sample shows evidence of a broad \siv\ trough without any significant trough present from the excited state line, which implies that this outflow could be at a distance of several kpc. Given the fractions of \civ\ and \siv, we establish firm limits on the global covering factor on \siv\ that ranges from 2.8\% to 21\% (allowing for the k-correction). Comparison of the expected optical depth for these ions with their detected percentage suggests that these species arise from common outflows with a covering factor closer to the latter. | 12 | 6 | 1206.7026 |
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1206 | 1206.0095_arXiv.txt | The \textit{Sun Watcher using Active Pixel system detector and Image Processing}(SWAP) on board the \textit{PRoject for OnBoard Autonomy\todash 2} (PROBA\todash 2) spacecraft provides images of the solar corona in EUV channel centered at 174~\AA. These data, together with \textit{Atmospheric Imaging Assembly} (AIA) and the \textit{Helioseismic and Magnetic Imager} (HMI) on board \textit{Solar Dynamics Observatory} (SDO), are used to study the dynamics of coronal bright points. The evolution of the magnetic polarities and associated changes in morphology are studied using magnetograms and multi-wavelength imaging. The morphology of the bright points seen in low-resolution SWAP images and high-resolution AIA images show different structures, whereas the intensity variations with time show similar trends in both SWAP 174 and AIA 171 channels. We observe that bright points are seen in EUV channels corresponding to a magnetic-flux of the order of $10^{18}$~Mx. We find that there exists a good correlation between total emission from the bright point in several UV\todash EUV channels and total unsigned photospheric magnetic flux above certain thresholds. The bright points also show periodic brightenings and we have attempted to find the oscillation periods in bright points and their connection to magnetic flux changes. The observed periods are generally long (10\todash 25~minutes) and there is an indication that the intensity oscillations may be generated by repeated magnetic reconnection. | \label{S-Introduction} X-ray bright points (XBPs) are prominent dynamical features in quiet-Sun and coronal hole regions. These features were first seen in rocket X-ray telescope images and were reported as small X-ray emitters with a spatial size of less than 60~arcseconds and lifetime ranging from a few hours to a few days \cite{1973ApJ...185L..47V}. \inlinecite{1974ApJ...189L..93G} determined the properties of XBPs from a study of \textit{Skylab} X-ray images. They found the lifetime of the bright points to be about eight hours in X-ray images. Bright points are also observed in EUV images \cite{1981SoPh...69...77H} where the emission has an average lifetime of approximately 20~hours as reported by \inlinecite{2001SoPh..198..347Z}. \inlinecite{1976SoPh...49...79G} estimated the average number of bright points emerging each day to be about 1500. They have also observed a class of long-lived bright points that exists near the Equator, with an average lifetime of about 36~hours. Bright points are always found to be associated with small, opposite-polarity poles in photospheric magnetograms, which have typical total flux of $10^{19}--10^{20}$~Mx \cite{1976SoPh...50..311G}. XBPs are likely signatures of small loops that connect the opposite polarities of some small-scale bipoles. It is estimated that one third of the bright points lie over ephemeral regions, which are newly emerging regions of magnetic flux, whereas the remaining two thirds lie above canceling magnetic features, which consist of opposite polarity fragments that approach one another and disappear \cite{1985AuJPh..38..875H,1993SoPh..144...15W}. This process normally takes place at the network boundaries of super granular cells \cite{1990ApJ...352..333H,2003A&A...403..731M,1983SvAL....9..385E}. \inlinecite{2001ApJ...547.1100H} found that only magnetic flux above a threshold of $3\times 10^{18}$~Mx is associated with a noticeable brightening in the EIT Fe~{\sc xii} corona. Emerging bipoles typically appear as a brightening on \textit{Yohkoh/SXT} images for only a small percentage of bipoles and have shorter lifetime in X-rays than in EUV, as was found by \inlinecite{2001SoPh..198..347Z}. A simple explanation given by \inlinecite{2003A&A...398..775M} is that the appearance at a certain temperature depends strongly on the magnetic energy released and, therefore, on the bipole magnetic-field strength. \inlinecite{2004A&A...418..313U} showed that not only does the coronal emission rise due to an increase in emission area, but also because the BP emits more per unit area as the magnetic flux becomes stronger. Whenever there is enhanced emission it is always associated with enhanced magnetic flux. They also found that there is a one-to-one correlation between the magnetic flux in the bipolarity and the EIT Fe~{\sc xii} coronal emission for the two BPs, both in the growing and decaying phase. The exception is that the brightest moments of the BP's lifetime, when there is strong increase in emission, are not accompanied by a comparable increase in the magnetic flux. The converging-flux model proposed by \inlinecite{1994ApJ...427..459P} describes how the approach of two opposite polarities creates an X-point that rises into the corona and produces a XBP by coronal reconnection \cite{1994SoPh..151...57P,1998ApJ...507..433L}. The model proposes a three-phase evolution: pre-interaction (approach), interaction, and finally cancellation. Several observational questions are raised concerning the evolution of these features. It would be interesting to check the timing of the disappearance of the emission at different temperatures using new observations. Another set of important characteristics of BP dynamics are oscillations in intensity. Several studies in EUV and soft X-ray spectral lines have reported a wide range of periodicities \cite{1979SoPh...63..119S,1979SoPh...63..113N,1992PASJ...44L.161S,2011MNRAS.415.1419K,2008A&A...489..741T}. \inlinecite{1979SoPh...63..119S} found that the constituent loops could evolve on a time scale of $\approx$ six minutes. \inlinecite{1981SoPh...69...77H} and \inlinecite{1990ApJ...352..333H}, using \textit{Skylab}, showed that BPs exhibit large variations in the emission of chromospheric, transition region, and coronal lines, and no regular periodicity or obvious correlation between the different temperatures was found. It is claimed that acoustic waves, which leak through the magnetic field lines of the solar atmosphere, can be converted into magnetoacoustic waves in the region where the plasma $\beta$ tends to unity and can reach the upper solar atmosphere \cite{2003ApJ...599..626B,2008AnGeo..26.2983K,2010MNRAS.405.2317S}. It is also reported that various intensity oscillations may also be generated by repeated magnetic reconnection \cite{2004PhDT.........1U,2005A&A...442.1087P,2006A&A...446..327D,2008A&A...489..741T}. Oscillations and the overall temporal evolution of BPs are especially important considerations for determining BPs' overall contribution to the coronal heating problem. In this article we discuss the evolution and dynamics of EUV bright points. We study the formation, lifetimes, temporal evolution and association with the photospheric magnetic field of several BPs in SWAP, and AIA images. | % \label{S-conclusions} Bright points are found to be always associated with opposite-polarity poles in photospheric magnetograms \cite{1976SoPh...49...79G}. Nevertheless, we see a lot of activity in HMI magnetograms corresponding to the location of BP1. We see multiple connectivities in BP1. From the formation stage of the bright point, and throughout its entire lifetime, there is continuous emergence of new magnetic flux and there is good overall correlation between the total magnetic flux and emission from the bright point in several UV and EUV channels. This suggests that the emerging flux interacts with the existing fields, which results in reconnection at the site of bright point. This, we conjuncture, leads to a series of continuous periodic brightenings as seen in different EUV lines. Within the small region around BP1's location, we see at least three major connectivities in the formation stage. We also observe that, in high-resolution images, the bright point looks like a miniature active region with multiple connectivities and with several loop structures. \inlinecite{2001ApJ...547.1100H} studied quiet-Sun loops, from EIT and MDI data. They have given a threshold of $3 \times 10^{18} $~Mx of magnetic flux to be associated with brightening in Fe~{\sc xii} corona. This number depends on the loop foot-point distance we see in EUV lines and the distance of the corresponding bipoles in the photosphere. In the case of bright points, which are small loops compared to typical quiet-Sun loops, this number differs. It also depends on which of the EUV lines we see brightenings in. In the case of BP1 we see brightenings in different channels of AIA and SWAP corresponding to a few times $10^{18}$~Mx. This supports the proposition discussed by \inlinecite{2003A&A...398..775M} that appearance of brightening at a certain temperature depends strongly on the magnetic energy released and therefore, on the magnetic-field strength. Combining our observational results for BP1 with bright-point observations by \inlinecite{2003A&A...398..775M} and observations of quiet coronal loops by \inlinecite{2001ApJ...547.1100H} one can conclude that size of the loops is also important as far as the brightening is concerned. Several authors \cite{1999ApJ...510L..73P,2001ApJ...547.1100H,2004A&A...418..313U,2008A&A...492..575P} reported a positive correlation between the EUV and X-ray emission with the underlying photospheric magnetic flux. \inlinecite{1999ApJ...510L..73P} observed that the photospheric magnetic flux is highly temporally correlated with X-ray and EUV emissions. \inlinecite{2004A&A...418..313U} found that the unsigned magnetic flux correlates well not only with the total EIT flux but also with EIT flux per pixel. They suggested that, as the magnetic flux becomes stronger, coronal emission from the BP increases not just because the emission area increases, but also because the intrinsic emission of the BP itself increases. For BP1, although we studied only a part of its lifetime, there exists a good correlation between the total magnetic flux and the total emission in different temperature channels (Figure~\ref{F-lcplot}a). But when the average emission per pixel is considered (Figure~\ref{F-lcplot}b), the correlation is not so good in the high-temperature channels, particularly during the initial phase of the bright point. However, the correlation is good for the entire duration in the channels that represent transition-region temperatures or below. So, the relationship between the intrinsic emission of the BP with increasing in magnetic flux indeed requires further confirmation. Nevertheless, the good correlation observed between total emission from the BP in six different channels, representing different temperature layers above the photosphere, indicates strongly that this BP is fed by magnetic reconnection; supporting the prevailing idea. \inlinecite{2008A&A...492..575P} have studied the structure and dynamics of a bright point from \textit{Hinode}, SOHO, and TRACE data. They found a positive correlation between weaker magnetic flux associated with one of the footpoints in a bipole and the total XRT brightening. In the case of BP1, negative flux is observed to be much stronger than the positive flux, but the emission is correlated well with the total unsigned flux, rather than with the fluxes from individual polarities. This result will have implications for both coronal heating and filling factor. In particular, we must point out here that we have selected only a few bright points to study the dynamical properties of CBPs; in particular we have focused our attention on one representative BP1 to show the general behavior of the BPs. A detailed statistical work will be essential in order to address whether or not the CBPs will be able to supply substantial energy flux to the energy budget of coronal heating. Finally we attempted to study whether or not the periodic brightenings of the CBPs can be used for coronal seismology. The observed periods are generally longer than those observed previously (10\todash 25~minutes, See Tables~1 and 2). There are some indications that the period ratio P1/P2 changes during the lifetime of the bright point, which may suggest a scenario of either magnetic-field divergence or density stratification as claimed by \inlinecite{2011MNRAS.415.1419K}. But we also must point out that we have not yet found a convincing theory that the period ratio can indeed be used as a tool for inferring the driver of such periodic brightenings. However, a good correlation between magnetic flux associated with the foot points and the intensity brightenings provides support for a repeated reconnection scenario. Further statistical work on this is required to come to a conclusion, which we intend to do in the future. \begin{acks} We thank the referee for their valuable comments. The AIA data used here is courtesy of SDO (NASA) and AIA consortium. DB and SKP thank the Guest Investigator Program of the PROBA2 mission, which supported this study. We would also like to thank the SWAP team, in particular D. Berghmans and Anik de Groof for their help at different stages of reduction of SWAP data. SWAP is a project of the Centre Spatial de Liege and the Royal Observatory of Belgium funded by the Belgian Federal Science Policy Office (BELSPO). \end{acks} | 12 | 6 | 1206.0095 |
1206 | 1206.5679_arXiv.txt | {The Large Magellanic Cloud (LMC) is rich in supernova remnants (SNRs) which can be investigated in detail with radio, optical and X-ray observations. \SNR\ is an X-ray and radio-bright remnant in the LMC, within which previous studies revealed the presence of a pulsar wind nebula (PWN), making it one of the most interesting SNRs in the Local Group of galaxies.} {We study the emission of \SNR\ to improve our understanding of its morphology, spectrum, and thus the emission mechanisms in the shell and the PWN of the remnant.} {We obtained new radio data with the Australia Telescope Compact Array and analysed archival \xmm\ observations of \SNR. We studied the morphology of \SNR\ from radio, optical and X-ray images and investigated the energy spectra in the different parts of the remnant.} {Our radio results confirm that this LMC SNR hosts a typical PWN. The prominent central core of the PWN exhibits a radio spectral index $\alpha_{\rm Core}$ of --0.04$\pm$0.04, while in the rest of the SNR shell the spectral slope is somewhat steeper with $\alpha_{\rm Shell}$ = --0.43$\pm$0.01. We detect regions with a mean polarisation of $P\cong$ (12$\pm$4)\% at 6~cm and (9$\pm$2)\% at 3~cm. The full remnant is of roughly circular shape with dimensions of (31$\pm$1)~pc $\times$ (29$\pm$1)~pc. The spectral analysis of the \xmm\ EPIC and RGS spectra allowed us to derive physical parameters for the SNR. Somewhat depending on the spectral model, we obtain for the remnant a shock temperature of around 0.2~keV and estimate the dynamical age to 12\,000-15\,000 years. Using a Sedov model we further derive an electron density in the X-ray emitting material of 1.56\,cm$^{-3}$, typical for LMC remnants, a large swept-up mass of 830~\msun, and an explosion energy of 7.6$\times 10^{50}$~erg. These parameters indicate a well evolved SNR with an X-ray spectrum dominated by emission from the swept-up material.} {} | \label{intro} The Magellanic Clouds (MCs) are considered to be one of the most ideal environments when it comes to the investigation of various supernova remnant (SNR) types and their different evolutionary stages. While their relatively small distance is very favourable for detailed studies, the MCs are also located in one of the coldest parts of the radio sky, allowing us to detect and investigate radio emission with little disturbing Galactic foreground radiation \citep{1991A&A...252..475H}. As they are located outside of the Galactic plane, the influence of dust, gas and stars is small, reducing the absorption of soft X-rays. Particularly, the Large Magellanic Cloud (LMC) at a distance of 50~kpc \citep{2008MNRAS.390.1762D}, allows for detailed analysis of the energetics of various types of remnants. In the radio-continuum regime SNRs are well-known for their strong and predominantly non-thermal radio emission, which is characterised by a typical spectral index of $\alpha\sim-0.5$ (as defined by $S\propto\nu^\alpha$). However, this value shows a significant scatter due to the wide variety of SNR types and this variance can be used as an age indicator for the SNR \citep{1998A&AS..130..421F}. SNRs have a significant influence on the structure of the interstellar medium (ISM). The appearance of spherically symmetric shell-like structures is very often perturbed by interaction with a non-homogeneous structure of the ISM. SNRs influence the behaviour, structure and evolution of the ISM. In turn, the evolution of SNRs is dependent on the environment in which they reside. Here, we report on new radio-continuum and archival X-ray observations of \SNR, one of five SNRs in the LMC (B0540-693, N157B, B0532-710, DEM\,L241 and \SNR) with a known or candidate pulsar wind nebula (PWN) inside. \SNR\ was initially classified as an SNR based on the \einstein\ X-Ray survey by \citet[][source 1 in their catalogue]{1981ApJ...248..925L} and \citet[][their source 2]{1991ApJ...374..475W}. \citet{1983ApJS...51..345M} catalogued this object based on studies of optical and Molonglo Synthesis Telescope (MOST) survey data, and reported an optical size of $140\arcsec\times131\arcsec$. \SNR\ is also listed in the 408~MHz MC4 catalogue of \citet{1976AuJPA..40....1C} as a distinctive point-like radio source, for which \citet{1983ApJS...51..345M} later re-measured a flux density of 350~mJy. \citet{1984AuJPh..37..321M} detected this source with specific MOST pointings and reported a rather flat spectral index of \mbox{$\alpha$=--0.38.} \ROSAT\ detected X-ray emission from \SNR\ during the all-sky survey \citep{1998A&AS..127..119F} and in deeper pointed observations using the PSPC \citep[source 670 in the catalogue of][]{1999A&AS..139..277H} and HRI detector \citep[source 8 in][]{2000A&AS..143..391S}. \citet{1998A&AS..130..421F} added further confirmation from radio data over a wide frequency range. \citet{1999ApJS..123..467W} classified it as an SNR with a ``diffuse face". \citet{1998ApJ...505..732H} presented X-ray spectra based on \asca\ observations. Their spectral modelling did not take into account the hard X-ray emission of the PWN, which was discovered later, and therefore overestimated its temperature. The most detailed study of \SNR\ was performed by \citet{2003ApJ...594L.111G} based on high resolution radio data obtained in 2002 with the Australia Telescope Compact Array (ATCA) and \cxo\ X-ray data from 2001. This lead to the discovery of the PWN inside the SNR and an age estimate of around 13\,000 years. \citet{2006ApJS..165..480B} reported no detection at far ultraviolet wavelengths based on data from the FUSE (Far Ultraviolet Spectroscopic Explorer) satellite. \citet{2006ApJ...652L..33W} detected this SNR in their Spitzer IR surveys as the object with the highest 70\,$\mu$m to 24~$\mu$m flux ratio of any SNR. However, due to high background emission at 70\,$\mu$m, in particular in the south western region of the remnant, they could only investigate the northern rim of the SNR shell. No radio pulsar within the area of \SNR\ was found in the systematic search within the LMC performed by \citet{2006ApJ...649..235M}. Also \citet{2008MNRAS.383.1175P} did not detect \SNR\ in their optical spectroscopic survey of LMC SNRs. Finally, \citet{2009ApJ...706L.106L} classified this LMC object based on its circular morphology as a core-collapse SNR. | \label{discussion} The \SNR\ in the LMC shows a prominent central core in radio and X-ray images, which is interpreted as a PWN \citep{2003ApJ...594L.111G}. The PWN is surrounded by a slightly elongated shell centred at RA(J2000) = 4\rahour53\ramin37\fs2, DEC(J2000) = --68\degr29\arcmin30\arcsec\ for which we derive a size of (128\arcsec$\pm$4\arcsec) $\times$ (120\arcsec$\pm$4\arcsec) (corresponding to (31$\pm$1~pc) $\times$ (29$\pm$1~pc) at a distance of 50~kpc) from the 13~cm radio data (obtained at position angles of 20\degr\ and 110\degr, running from north to east). The shell radio emission shows brightening along its southern rim. Our size estimate is in agreement with the diameter previously reported from radio data by \citet{2003ApJ...594L.111G}. From the intensity profile of the EPIC data we estimate a size of $\sim$140\arcsec\ (roughly in east - west direction, at 10\% of maximum intensity), which corresponds to $\sim$130\arcsec\ corrected for the angular resolution of the telescope. For a better estimate we re-analysed the \cxo\ data from 2001 presented by \citep{2003ApJ...594L.111G}. From the \cxo\ image we measure 128\arcsec$\pm$4\arcsec (31$\pm$1~pc) as the largest extent, which is consistent with the radio extent (Fig.~\ref{chandra}). So in the following considerations we assume a radius of 15.5~pc. \begin{figure}[h] \begin{center} \resizebox{0.98\hsize}{!}{\includegraphics[angle=-90,clip=]{0453mcels13cm.eps}} \caption{MCELS composite optical image \textrm{(RGB=H$_\alpha$,[S\textsc{ii}],[O\textsc{iii}])} of \SNR\ overlaid with 13~cm contours. The contours are 0.12, 0.4, 1, 4 and 8~mJy/beam.} \label{mcels} \end{center} \end{figure} To estimate the age of the SNR from the vpshock model we used the relation t$_{\rm y}$\,=\,3.8\expo{2}\,R$_{\rm pc}$(kT)$_{\rm keV}^{-1/2}$ from \citet{2005ChJAA...5..165X} and assume the temperature derived from X-ray spectral modelling (0.25~keV) and the radius of the SNR, which results in $\sim$11.8~kyr, consistent with the age inferred by \citet{2003ApJ...594L.111G} from a preliminary X-ray spectral modelling. Also using a similar Sedov model and following \citet{2011A&A...530A.132O} who applied it to the spectrum of the SMC SNR IKT\,16, we derive -- with a shock temperature of 0.17 keV, and the normalisation of the Sedov model fit of 3.74\expo{-2} cm$^{-5}$ -- the following physical parameters for \SNR: electron density in the X-ray emitting material of 1.56\,cm$^{-3}$, dynamical age of 15.2~kyr, swept-up mass of 830\,\msun\ and an explosion energy of 7.6\expo{50}~erg. Our estimate for the density is higher than that derived by \citet{1998ApJ...505..732H} from \asca\ observations (without knowledge of the PWN contributing to the X-ray spectrum of the remnant) and consistent with the average value of the other LMC SNRs in that work. This directly reflects the higher density of the ISM in the LMC in comparison to the SMC \citep{2004A&A...421.1031V}. \begin{figure}[t] \begin{center} \resizebox{0.98\hsize}{!}{\includegraphics[angle=-90,clip=]{color0453_Chandra_vs_13cm.eps}} \caption{\cxo\ colour image of \SNR\ in the 0.3$-$2.0 keV (red) and 2.0$-$4.0 keV (green) energy bands overlaid with 13~cm contours. The contours are 0.12, 0.4, 1 and 4~mJy/beam and the \cxo\ image was smoothed to the same resolution as the radio image.} \label{chandra} \end{center} \end{figure} The shock temperature (and electron temperature in our Sedov modelling) is very low compared to other remnants in the MCs. The whole SMC sample modelled by \citet{2004A&A...421.1031V} and \citet{2008A&A...485...63F} in a similar way exhibits temperatures between 0.26~keV and 1.8~keV. The low temperature leads to a relatively high dynamical age for \SNR. The high swept-up mass is also consistent with an SNR well evolved into its Sedov phase. As discussed in \citet{2004A&A...421.1031V} the higher ISM density and higher abundances in the LMC compared to the SMC can lead to a faster evolution to the radiative cooling stage. The abundances derived from our spectral analysis exhibit some model dependences, but are generally consistent or somewhat lower than average abundances in the ISM of the LMC. The only exception is Si which is overabundant with respect to the other elements. This suggests that the X-ray spectrum is dominated by emission from swept-up ISM material, which makes it difficult to draw conclusions on the type of supernova explosion from abundance measurements. In the case of \SNR\ the association of the PWN with the remnant and its morphology favour a core-collapse supernova \citep{2009ApJ...706L.106L}. The radio spectral index of $\alpha_{\rm Shell}=-0.43$, confirms the non-thermal nature of the SNR shell emission in the radio band, while the flat spectral index for the core of $\alpha_{\rm Core}=-0.04$ is typical for a PWN. The overall spectral index of $\alpha=-0.39$ is slightly ``flatter'' in comparison with typical values of $-0.5$ for SNRs \citep{1985ApJS...58..197M,1998A&AS..130..421F}. The radio spectra definitely confirm the PWN nature of the central object. At higher frequencies (Fig.~\ref{specidx}) the flux density decreases as expected whereas at lower frequencies non-thermal radiation of the shell dominates. Overall, the radio-continuum properties of \SNR\ are very similar to SNR~B0540-693, the most prominent SNR with a PWN in the LMC \citep{1993ApJ...411..756M}. \citet{1983ApJS...51..345M} found distinctive optical connections to X-ray and radio features of this SNR. As can be seen in the MCELS images (Fig.~\ref{mcels}), the \OIII\ emission dominates in the outer parts of the remnant. This could indicate an oxygen-rich type of SNR and suggests a type~Ib SN event -- the explosion of a massive O, B, or WR star \citep{2005MNRAS.360...76A}. We also note that prominent H$\alpha$ emission is visible in the south-west, % which might suggest that higher density ISM material is causing the brightening in the radio shell emission at the outer rim of the SNR in this direction. While no molecular clouds are reported in this region \citep{2008ApJS..178...56F,2009ApJS..184....1K}, we do note that \SNR\ lies in a region of low \ion{H}{i} column density with the rim density highest in the south west \citep{1998ApJ...503..674K,1999AJ....118.2797K}. The radio spectral index (Fig.~\ref{specmap}) in this south-west region indicates non-thermal emission $(\alpha \sim -0.4)$. Careful comparison (after correcting for the rotated presentation) of the 24 \& 70 $\mu$m infra-red data shows no correspondence between the enhanced background reported by \citet{2006ApJ...652L..33W} and \SNR. In contrast however, the \cxo\ image reveals only weak X-ray emission at the location of the bright radio feature. This could be explained by the higher density material being located in front of the rim of the SNR, suppressing the X-rays, or, together with the high polarisation of the 6\,cm radio emission in this area, might favour a locally increased magnetic field strength enhancing the non-thermal radio emission. Future spatially resolved X-ray spectroscopy should be able to resolve this question. | 12 | 6 | 1206.5679 |
1206 | 1206.5347_arXiv.txt | {} {We present first results of our efforts to re-analyze the \emph{Kepler} photometric dataset, searching for planetary transits using an alternative processing pipeline to the one used by the \emph{Kepler} Mission} {The SARS pipeline was tried and tested extensively by processing all available \emph{CoRoT} mission data. For this first paper of the series we used this pipeline to search for (additional) planetary transits only in a small subset of stars - the \emph{Kepler} objects of interest (KOIs), which are already known to include at least one promising planet candidate.} {Although less than 1\% of the \emph{Kepler} dataset are KOIs we are able to significantly update the overall statistics of planetary multiplicity: we find \nNew new transit signals on \nSystems systems on these light curves (LCs) only, nearly doubling the number of transit signals in these systems. \nWereSingleUpper of the systems were singly-transiting systems that are now multiply-transiting. This significantly reduces the chances of false positive in them. Notable among the new discoveries are KOI 435 as a new six-candidate system (of which kind only Kepler-11 was known before), KOI 277 (which includes two candidates in a 6:7 period commensurability that has anti-correlated transit timing variations) -- all but validating the system, KOIs 719, 1574, and 1871 that have small planet candidates ($1.15, 2.05$ and $1.71 R_\oplus$) in the habitable zone of their host star, and KOI 1843 that exhibits the shortest period (4.25hr) and among the smallest (0.63 $R_\oplus$) of all planet candidates. We are also able to reject \nReject KOIs as eclipsing binaries based on photometry alone, update the ephemeris for \nEphem KOIs and otherwise discuss a number of other objects, Which brings the total of new signals and revised KOIs in this study to more than one hundred. Interestingly a large fraction, about $\sim 1/3$, of the newly detected candidates participate in period commensurabilities. Finally, we discuss the possible overestimation of parameter errors in the current list of KOIs and point out apparent problems in at least two of the parameters.} {Our results strengthen previous analyses of the multi-transiting ensemble, and again highlight the great importance of this dataset. Nevertheless, we conclude that despite the phenomenal success of the \emph{Kepler} mission, parallel analysis of the data by multiple teams is required to make full use of the data.} | The \emph{Kepler} mission goal is to determine the frequency of Earth-like planets around Sun-like stars, and in a wider sense the distribution of planets of all types around a variety of host stars (Borucki \etal 2010 and references therein). \emph{Kepler} has already produced a large number of milestone results (e.g. see introduction of Batalha \etal 2012 (hereafter B12) for a complete list), its importance to exoplanet studies is immense now and is expected to be even larger in the future. The previous list of over 1200 Kepler candidates (Borucki \etal 2011b) was severely incomplete. Indeed, the \emph{Kepler} team recently almost doubled the number of candidates using improved detection tools (and only slightly more data) (B12). Even citizen scientists who just browsed the huge database were able to find several such missed signals (Fischer \etal 2012, Lintott \etal 2012) as an alternative approach to finding more candidates. We therefore checked the completeness of the updated \emph{Kepler} candidate list and report here that we are able to find a significant number of very good candidates that were missed by the \emph{Kepler} team. The \emph{Kepler} pipeline meticulously processes the raw pixels transmitted from the spacecraft to produce the outstanding results achieved so far. This pipeline is progressively perfected, and together with the accumulating data has produced lists of growing length from 705 (Borucki \etal 2011a) to 1235 (Borucki \etal 2011b) to 2321 (B12) transiting-planet candidates. Still, the analysis of the raw data and subsequent transit searches that is done by the \emph{Kepler} team is to date almost the only analysis of this dataset (the very recent work by Huang \etal 2012, hereafter H12, is the only exception -- more below). This is in sharp contrast to the \emph{CoRoT} space mission, where the calibrated light curves (LCs) are distributed to a number (up to 10) of different teams and each team tries its own tool set to detect the most and best candidate signals, which are then individually approved or rejected in a common discussion. This approach allows a more complete surveying of the dataset for transiting planets. In this paper series we apply this approach to the \emph{Kepler} dataset and re-analyze it entirely with our own tool set. So far, very few attempts to do this are known: there is the Planet Hunters crowd-sourcing search (Fischer \etal 2012, Lintott \etal 2012) (which did not re-process the data) that found four, and the HATNet team (H12) that found 16 good periodic signals that were not identified by the \emph{Kepler} pipeline using Q0-Q6 data. In this first paper of the series we use a little-changed \emph{CoRoT}-oriented pipeline to remove systematic effects from the whole Q0-Q6 \emph{Kepler} data set. We then search for additional transit-like signals only on a subset of less than 1\% of all cleaned LCs - the \emph{Kepler} KOIs (Kepler Objects of Interest). All KOIs have already at least one promising transit candidate in them and passed a battery of tests for false positives (B12). Furthermore, the detection of a reliable second (or third, etc.) signal in the same LC dramatically reduces the remaining chances of false positives so that nearly all multi-transiting LCs are indeed of planetary origin (Lissauer \etal 2012, Fabrycky \etal 2012a). The true distribution of single- versus multi- transiting planets systems in itself is also of interest to understand the planet formation theory (Figueira \etal 2012). Other papers in this series are expected present the application of the same analysis to the entire \emph{Kepler} dataset (which is $>100$ times larger than just the KOIs) or to highlight particular systems. The paper is organized in the following way: in \S \ref{SARSprocessing} we describe our pre-processing and detection pipeline and also discuss the fitting of multi-transiting systems. In \S \ref{NewCandidates} we present our new detections, in \S \ref{Notes} we discuss other KOIs that are determined to be eclipsing binaries (hereafter EBs) or are otherwise revised or noted, and we conclude in \S \ref{Discuss}. Appendix \ref{appA} includes the remaining graphical fits of systems with new detected candidates that were not presented earlier, and Appendix \ref{appB} has additional graphical content related to the other KOIs. | \label{Discuss} We presented the discovery of \nNew transit signals found on a subset of less than 1\% of the \emph{Kepler} dataset - the Kepler objects of interest (KOIs). These signals were not detected by the \emph{Kepler}'s processing pipeline in the publicly available Q0-Q6 data. We were able to achieve that with the SARS pipeline that is was optimized for \emph{CoRoT}, and not for \emph{Kepler}, and was left almost unchanged from it. Still, the very basic fact that this pipeline is simply different from the one used by \emph{Kepler} team means that its sensitivities and biases are also different. Furthermore, we showed evidence that for the specific goal of transit searches the SARS pipeline may even be superior to \emph{Kepler}'s PDC-MAP -- helping to understand the significant number of new detections in such a small subsample of LCs. It is interesting to note that in ststems previously known to be multi-transiting often the \emph{order} in which signals are detected was different in the SARS pipeline compared to the published order -- this is a hallmark of the different sensitivities of each pipeline. To boost the confidence level in the newly detected signals we employed a very large array of tests for candidate selection: to the regular binarity and contamination tests we added the ``half-half'' test for false-positives identification, required that signals are not effected by changing the long-term filter (and thus are even less likely to be an artifact of data processing), and sometimes confirmed that they are present in \emph{Kepler}'s PDC data as well. This suite of tests exceeds the one used by the \emph{Kepler} pipeline, helping to produce a clean catalog. We also showed evidence that the errors given in B12 for the $d$ and $r_p$ parameters are overestimated for a significant fraction of the KOIs, to the point where for about half the KOIs the value of $d$ is not significantly different from zero ($\frac{d}{\Delta d} <3$), which is highly unlikely to be true. The fact that additional transit candidates are found in the data of \nWereSingleLower systems that had only one previously known signal significantly boosts the chances that these systems are real planets since the false-alarm rate of multi-transiting planets is, intrinsically, very low. This is especially true for the KOI 277 system, which also exhibits anti-correlated TTVs in a dynamically stable configuration, all but validating the system. Also notable is the large fraction of planet candidates that are part of period commensurabilities, more than than presented in Fabrycky \etal (2012a). We do not have a good explanation for why this happened - but near-resonant systems are particularly useful because they may allow for the full confirmation of their planetary nature using photometry alone. We advocate the parallel processing of the full data by other teams. However, in the near future the publicly available \emph{Kepler} data will significantly increase -- and with it the associated workload related to analyzing it, both in computer time and human time. We believe that the sheer volume of the work means that re-analyzing the data will be increasingly more difficult, and thus increasingly less possible for small groups of researches - but we hope that the results of this work will encourage others to work in this direction. For instance, the \emph{Kepler} pipeline has five primary modules: CAL, PA, PDC, TPS, and DV for the calibration, photometry, removal of systematic effects, transit search, and signal validation, respectively. In this paper we presented our own versions of the PDC, TPS, and DV modules - and these changes were productive. We therefore believe that an independent analysis matching the CAL and PA modules (which remain ``unchallenged'' to date) may also be productive. Finally, since we were able to detect up to four additional transit signals in each of the presented systems, the chances of detecting more transiting systems -- and even more multi-transiting systems -- in the remaining $>99\%$ of the data are quite real. | 12 | 6 | 1206.5347 |
1206 | 1206.0562_arXiv.txt | We apply the integrated perturbation theory \cite{mat11} to evaluate the scale-dependent bias in the presence of primordial non-Gaussianity. The integrated perturbation theory is a general framework of nonlinear perturbation theory, in which a broad class of bias models can be incorporated into perturbative evaluations of biased power spectrum and higher-order polyspectra. Approximations such as the high-peak limit or the peak-background split are not necessary to derive the scale-dependent bias in this framework. Applying the halo approach, previously known formulas are re-derived as limiting cases of a general formula in this work, and it is implied that modifications should be made in general situations. Effects of redshift-space distortions are straightforwardly incorporated. It is found that the slope of the scale-dependent bias on large scales is determined only by the behavior of primordial bispectrum in the squeezed limit, and is not sensitive to bias models in general. It is the amplitude of scale-dependent bias that is sensitive to the bias models. The effects of redshift-space distortions turn out to be quite small for the monopole component of the power spectrum, while the quadrupole component is proportional to the monopole component on large scales, and thus also sensitive to the primordial non-Gaussianity. | Introduction } The primordial non-Gaussianity is a useful indicator in searching for the generation mechanism of density fluctuations in the universe. While the primordial non-Gaussianity is small in the simplest model of single-field slow-roll inflation, there are various other inflationary models which predict fairly large primordial non-Gaussianities (for review, see Refs.~\cite{BKMR04,che10}). Alternative models without inflation can also produce large non-Gaussianities (see, e.g., Ref.~\cite{leh10}). Accordingly, constraining or detecting the primordial non-Gaussianity has a substantial importance in studying very early stages of the universe. The large-scale structure (LSS) of the universe has been one of the most important ways of constraining cosmological models in general. In recent years, it was found that the primordial non-Gaussianity produces the scale-dependent bias in the LSS \cite{dal08,MV08,DS10a}. Although the scale-dependent bias from the primordial non-Gaussianity mainly appears on very large scales, the form of the scale dependence does not receive general relativistic corrections even on scales larger than the Hubble radius \cite{WS09}. The scale-dependent bias as a method to constrain the primordial non-Gaussianity has already been applied to observational data of galaxies and quasars \cite{slo08,xia10a,xia10b,xia11}. The constraints derived from the scale-dependent bias are competitive with the measurements in the cosmic microwave background (CMB). A hint of a positive value of local-type non-Gaussianity parameter $f_{\rm NL}$ was indicated by analyses of the scale-dependent bias \cite{xia11}, which could have profound implications for models of the early universe. However, more detailed analyses with large galaxy surveys are necessary to derive conclusive results. On the theoretical side, analytic expressions for the scale-dependent bias has been only approximately derived. So far, at least three kinds of derivation are known. One is based on the method of peak-background split \cite{dal08,AT08,slo08,GP10,SK10,DJS11a,DJS11b,SHMC11}, and the other is based on the statistics of high-threshold regions \cite{MV08,VM09,JK09,DJS11b}. It is also shown that the scale-dependent bias generally appears in phenomenological models of local bias \cite{mcd08,TKM08}. Because those theoretical derivations are approximate and not exact, they should have been compared with numerical simulations \cite{DSI09,gro09,PPH10,WVB10,WV11}. The scale-dependent bias is qualitatively confirmed by simulations, although the detailed amplitude of the analytic predictions needs to be modified \cite{DSI09,WV11,DJS11a}. The purpose of this paper is to provide more precise and more general analytic expressions for the scale-dependent bias in the presence of primordial non-Gaussianity. Evolutions of density fluctuations on sufficiently large scales are expected to be described by the nonlinear perturbation theory. However, a consistent inclusion of the general form of bias in the nonlinear perturbation theory had not been clear until recently. The integrated perturbation theory (iPT) \cite{mat11} is the formalism in which a broad class of biasing scheme can be consistently included on a general ground. The local Eulerian biasing scheme has been frequently adopted in attempts of incorporating the bias into the nonlinear perturbation theory. However, the local Eulerian bias is not consistent with the nonlinear dynamics in general \cite{mat11,CSS12}, because the nonlinear evolutions are nonlocal phenomena. In the formalism of iPT, generally nonlocal biasing either in Eulerian and Lagrangian spaces can be consistently included in the nonlinear perturbation theory. The effects of primordial non-Gaussianity and redshift-space distortions are naturally incorporated into the formalism. Consequently, it is a straightforward application of iPT to making predictions of scale-dependent bias in the presence of primordial non-Gaussianity. In this paper, we first present the most general prediction of iPT for the scale-dependent bias, which can be applied to almost any type of primordial non-Gaussianity and to almost any model of bias, as long as we consider the regime where the perturbation theory applies. We then consider popular models of primordial non-Gaussianity, i.e., local-, equilateral-, folded-, and orthogonal-type non-Gaussianities. Asymptotic behaviors of scale-dependent bias on large scales, which have been derived in limited cases in the literature, are re-derived in general situations without resorting to specific forms of bias. When the halo model of bias is adopted, the previously known formulas of scale-dependent bias are re-derived by taking appropriate limits of our general formula. In course of derivation, nonlocality of bias turns out to be important. We also show that previous formulas derived from the peak-background split are only consistent with the Press-Schechter mass function. When the mass function deviates from the Press-Schechter form, our general formula predicts that those previous formulas of scale-dependent bias should be corrected. The main purpose of this paper is to show how the iPT can be applied to making theoretical predictions of the scale-dependent bias, and to give lowest-order calculations with primordial bispectra, when the bias is given by a simple, nonlocal model of halo bias. With the iPT formalism, predicting the scale-dependent bias is straightforward once the bias model is given. Theoretical uncertainties in predicting the scale-dependent bias are reduced to those of the bias model. Accurate modeling of biasing is actively studied in recent years. Once we have an accurate model of galaxy bias which is generally nonlocal, the iPT immediately gives predictions of scale-dependent bias with least approximations. This paper is organized as follows. In Sections \ref{sec:SDBias}--\ref{sec:BiasFn}, analytic derivations of scale-dependent bias in real space are presented in order. The general formula of the biased power spectrum in real space with primordial bispectrum is derived by the lowest-order iPT in Sec.~\ref{sec:SDBias}. Large-scale limits of the scale-dependent bias in concrete models of primordial non-Gaussianity are generally investigated in Sec.~\ref{sec:PNGmodels} without assuming the models of bias. In Sec.~\ref{sec:BiasFn}, we consider the shapes of renormalized bias functions which are needed in iPT. We generalize the previous results to include the effects of the smoothing function in the halo model of bias. Numerical evaluations of the analytic formula are presented and compared with previous predictions in Sec.~\ref{sec:Numerical}. In Sec.~\ref{sec:SDBiasR}, we generalize our formula to include the redshift-space distortions. In Sec.~\ref{sec:concl}, we summarize our results. In plotting the figures of this paper, we adopt a cosmological model of flat curvature with parameters $\varOmega_{\rm m0} = 0.275$, $\varOmega_{\Lambda 0} = 1-\varOmega_{\rm m0} = 0.725$, $\varOmega_{\rm b0} = \varOmega_{\rm m0} - \varOmega_{\rm c0} = 0.046$, $n_{\rm s} = 0.96$, $\sigma_8 = 0.8$, $H_0 = 70\,{\rm km/s/Mpc}$, where $\varOmega_{\rm m0}$ is the matter density parameter, $\varOmega_{\Lambda 0}$ is the cosmological constant parameter, $\varOmega_{\rm b0}$ is the baryon density parameter, $\varOmega_{\rm c0}$ is the cold dark matter density parameter, $n_{\rm s}$ and $\sigma_8$ are respectively the spectral index and the amplitude of primordial density fluctuations, and $H_0$ is the Hubble's constant. | } The iPT is a general framework of the perturbation theory in the presence of bias. In this paper, we first apply this framework to deriving the relation between the scale-dependent bias and the primordial non-Gaussianity. Approximations such as the peak-background split and the high-peak limit, which are usually adopted in the literature to estimate the scale-dependent bias, are not required. The redshift-space distortions of the scale-dependent bias are also evaluated. Thus, in this paper, we have derived the most general formula so far of the scale-dependent bias with primordial non-Gaussianity in the literature. For the scale-dependent bias in real space, the most fundamental equation in this paper is provided by Eq.~(\ref{eq:2-13}), where $Q_n(k)$ is linearly dependent on the primordial bispectrum. We find that the slope of the scale-dependent bias in the large-scale limit is determined only by primordial bispectra in the squeezed limit, and is independent on detailed models of bias. This property explains the fact that different models of bias have predicted the same slope of the scale-dependent bias in the literature. We derive the shape of renormalized bias functions, generalizing the concept of simple Press-Schechter approach. The general expression of renormalized bias functions in this approach is given by Eq.~(\ref{eq:4-26}). In the case of universal mass function, the renormalized bias functions are given by Eq.~(\ref{eq:4-32}), or equivalently Eq.~(\ref{eq:4-33}). The previously known results in the approximation of peak-background split are reproduced when the PS mass function is assumed in our results [Eq.~(\ref{eq:4-53})]. When the mass function deviates from the PS form, our results suggest that the form of scale-dependent bias should be corrected. The general formula of scale-dependent bias on large scales is given by Eq.~(\ref{eq:4-54}). This equation is one of the main outcomes in this paper. Most of the previous results regarding the scale-dependent bias from the primordial non-Gaussianity are derived as special cases of this equation. The evaluations of the redshift-space distortions in the scale-dependent bias are straightforward in the framework of iPT. The result is given by Eq.~(\ref{eq:6-9}), or in terms of multipole coefficients, Eqs.~(\ref{eq:6-14a})--(\ref{eq:6-14d}). On large scales, however, dominant terms in these equations are simply given by Eqs.~(\ref{eq:6-15a}) and (\ref{eq:6-15b}). When the bias is large enough, the redshift-space distortions do not affect the scale-dependence of bias much [Eqs.~(\ref{eq:6-16a}) and (\ref{eq:6-16b})]. Even when the bias is not large enough, the redshift-space distortions have little effects at least in the non-Gaussian models we have considered (Figs.~\ref{fig:rm114} and \ref{fig:rk005}). While highly biased objects have large amplitudes of power spectrum, the number of objects is small and the shot noise is large. Thus highly biased objects are not suitable for testing the primordial non-Gaussianity. On the other hand, the amplitude of power spectrum is small for less biased objects, and the clustering signals are small. Consequently, there should be an optimal objects with sufficiently large bias and sufficiently large numbers at the same time for realistically constraining the primordial non-Gaussianity by galaxy surveys. The high-peak limit or the peak-background split are not necessarily valid in some cases. The results of this paper provide the most accurate formula of the scale-dependent bias in the literature. They should be useful in theoretical investigations as well as in constraining the primordial non-Gaussianity with realistic galaxy surveys. Applications of the results in this paper, including the Fisher analysis of the future galaxy surveys, higher-order analyses of primordial non-Gaussianity, are now in progress. For more accurate modeling of the scale-dependent bias, it should be necessary to improve the nonlocal bias model beyond the simple halo approach. Investigations in this direction will be addressed in future work. | 12 | 6 | 1206.0562 |
1206 | 1206.5806_arXiv.txt | A small fraction of the atomic-cooling halos assembling at $z<15$ may form out of minihalos that never experienced any prior star formation, and could in principle host small galaxies of chemically unenriched stars. Since the prospects of detecting isolated population III stars appear bleak even with the upcoming {\it James Webb Space Telescope} (JWST), these population III galaxies may offer one of the best probes of population III stars in the foreseeable future. By projecting the results from population III galaxy simulations through cluster magnification maps, we predict the fluxes and surface number densities of pop III galaxy galaxies as a function of their typical star formation efficiency. We argue that a small number of lensed population III galaxies in principle could turn up at $z\approx 7$--10 in the ongoing {\it Hubble Space Telescope} survey CLASH, which covers a total of 25 low-redshift galaxy clusters. | The first generation of population III stars (hereafter pop III) is predicted to form in isolation or in small numbers within dark matter minihalos ($\sim 10^5$--$10^6\ M_\odot$) at $z\leq 60$ \citep[e.g.][]{Trenti&Stiavelli09}. However, even with the upcoming James Webb Space Telescope (JWST), the prospects of detecting such stars appear bleak \citep{Rydberg12}, unless they attain masses $\geq 10^2$--$10^3\ M_\odot$ due to a prolonged dark star phase \citep[e.g.][]{Freese10,Zackrisson10a,Zackrisson10b,Ilie12}. Larger numbers of pop III stars could in principle form at $z\leq 15$ within HI cooling halos ($\sim 10^7$--$10^8\ M_\odot$) that have remained chemically pristine \citep[e.g.][]{Johnson09,Stiavelli&Trenti10}, and such pop III ``galaxies'' (a.k.a. pop III star clusters) should be significantly easier to detect. Fig.~\ref{schematic} schematically illustrates how such objects might form. \begin{figure} \includegraphics[height=.4\textheight]{fig1} \caption{Schematic illustration of the formation of a pop III galaxy at high redshift. {\bf a)} A minihalo forms one or several pop III stars, which emit Lyman-Werner radiation (blue region) and inhibit star formation in adjacent minihalos. {\bf b)} The pop III stars explode as supernovae and eject metals (red region) into the surrounding medium. {\bf c)} The pop III minihalo merges with a number of nearby minihalos, reaches the HI cooling limit required for prolonged star formation and forms a ``first galaxy''. Due to metal enrichment provided by the pop III star/s in one of the progenitor minhalos, the galaxy becomes a pop II/I object. {\bf d)} A rare, previously sterilized minihalo that has remained chemically pristine reaches the HI cooling limit and forms a pop III galaxy.} \label{schematic} \end{figure} | Depending on their typical star formation efficiencies, lensed pop III galaxies may be within reach of JWST and possibly even the HST. Such objects may, under certain circumstances be identified in multiband photometric surveys because of their unusual broadband colours. A search for such objects in the HST/CLASH survey, targeting 25 foreground galaxy clusters, is currently underway. | 12 | 6 | 1206.5806 |
1206 | 1206.4031_arXiv.txt | We present the first time-series study of the ultra-faint dwarf galaxy Hercules. Using a variety of telescope/instrument facilities we secured about 50 $V$ and 80 $B$ epochs. These data allowed us to detect and characterize 10 pulsating variable stars in Hercules. Our final sample includes 6 fundamental-mode (ab-type) and 3 first overtone (c-type) RR Lyrae stars, and one Anomalous Cepheid. The average period of the ab-type RR Lyrae stars, $\langle$P$_{ab}\rangle $=0.68 d ($\sigma$ = 0.03 d), places Hercules in the Oosterhoff II group, as found for almost the totality of the ultra-faint dwarf galaxies investigated so far for variability. The RR Lyrae stars were used to obtain independent estimates of the metallicity, reddening and distance to Hercules, for which we find: $[Fe/H]=-2.30\pm0.15$ dex, $E(B-V)=0.09\pm0.02$ mag, and $(m-M)_0=20.6 \pm 0.1$ mag, in good agreement with the literature values. We have obtained a $V, B-V$ color-magnitude diagram (CMD) of Hercules that reaches $V \sim$ 25 mag and extends beyond the galaxy's half-light radius over a total area of $40' \times 36'$. The CMD and the RR Lyrae stars indicate the presence of a population as old and metal-poor as (at least) the Galactic globular clusters M68. | The discovery of a new class of faint dwarf satellites around the Milky Way (MW; see e.g., \citealt[and references therein]{bel06,Bel10}) and the Andromeda galaxies (M31; see e.g., \citealt[and references therein]{Richardson11}) has opened a new window for the study of the formation history of large spirals. The new systems show a number of remarkable differences with respect to the ``classical'' dwarf spheroidals (dSphs) surrounding the MW and M31: i) they have much lower surface brightnesses ($\mu_V \gtrsim 28$~mag), for which they were named ``ultra-faint'' dwarfs (UFDs); ii) they are very metal poor, with large dispersions and [Fe/H] values as low as $-4$ dex \citep[see][and references therein]{THT09}. Such extreme abundances are not observed among the classical dSphs where only a few stars with [Fe/H]$<-3.0$ have been detected \citep{Fr10} compared to the large number found in the MW halo; iii) they generally contain RR Lyrae stars that conform to the subdivision into Oosterhoff types I (Oo I) and II \citep[OoII;][]{Oo39}\footnote{The Galactic globular clusters can be divided into two different groups according to the mean period of their ab-type RR Lyrae stars: Oosterhoff type I (Oo I) clusters have $\langle P \rangle=0.55$ d, whereas type II (Oo II) clusters have $\langle P \rangle=0.65$ d} observed for field and cluster MW variables. So far, the only exception among the UFDs is Canes Venatici I \citep[CVn I,][]{Kuehn08}, the brightest of the MW UFDs, that, like the classical MW dSphs, has instead Oosterhoff-intermediate (Oo-Int) properties \citep[][and references therein]{Cat09,Clementini10}; and, finally: iv) the UFDs discovered so far outnumber by almost a factor of two the classical dSphs, thus partially reducing the so-called ``missing satellites problem'' \citep{Moore99,Kl99} affecting the $\Lambda$-cold-dark-matter ($\Lambda$CDM) scenario of galaxy formation. With their properties the UFDs are potentially much better analogues than the classical dSphs of the ``building blocks'' that contributed to the formation of the two large spirals in the Local Group. They have absolute luminosities similar to the globular clusters (GCs; $M_V\sim -7$ mag, on average) but they are much more spatially extended than GCs. With typical half-light radii of $r_h \ge$ 100 pc, in fact they equal in size the ``classical'' dSphs. The UFDs are found in groups on the sky \citep[see e.g. Fig. 1 of][]{Richardson11}, have small velocity dispersions, and high mass-to-light ratios. All UFDs host an ancient population around 10 Gyr old. They have GC-like color-magnitude diagrams (CMDs) resembling the CMDs of metal-poor Galactic GCs (GGCs) such as M92 (NGC~6341), M15 (NGC~7078) and M68 (NGC~4590). Some of the MW UFDs have a distorted shape due to the tidal interaction with the MW. We are carrying out an extensive observational campaign of the new MW and M31 UFDs to study structural parameters and stellar population properties, as well as the variable stars of these systems. We have already published results for Bootes I \citep[Boo I,][]{Dallora06}, CVn I \citep[][]{Kuehn08}, Canes Venatici II \citep[CVnII,][]{Greco08}, Coma \citep{Musella09}, Leo IV \citep{Moretti09}, and Ursa Major II \citep[UMaII][]{Dallora12} among the MW UFDs. With the only exception of CVn I, all these dwarfs contain RR Lyrae stars with pulsation periods suggesting an Oo II classification. However, only Boo I and CVn I contain sufficiently large numbers of variables to be safely classified into Oosterhoff types. The classification of the other UFDs is less certain given the small numbers of variables they contain. Nevertheless, their few variables clearly tend to have Oosterhoff type II properties. Thus, in terms of stellar metallicity and pulsation properties of the variable stars, systems similar to the UFDs, as they were at earlier times, could resemble the building blocks of the Galactic halo. In this paper we extend our analysis to the Hercules UFD galaxy (R.A.$=16^h 31^m 02.0^s$, Decl.$=12^{\circ}47^{\prime}29.6^{"}$, J2000.0). The galaxy was discovered by \citet{Bel07} from the analysis of SDSS data. The CMD of Hercules (based on follow-up Isaac Newton Telescope - INT- data) shows besides the red giant branch (RGB), also blue and red horizontal branches (HBs). \citet{Bel07} interpreted this evidence as a possible signature of multiple stellar populations present in the galaxy. A number of studies have been devoted to this galaxy after the discovery paper. CMDs of Hercules reaching well below the galaxy's main sequence turn-off were published by \citet[][in the $B$ and $V$ bands]{Col07} and \citet[][in the $g$ and $B$ bands]{Sand09}, based on very wide-field observations obtained with the red and blue channels of the Large Binocular Camera (LBC; \citealt{Giallongo08}) of the Large Binocular Telescope (LBT). According to the \citet[][]{Sand09} recovery of the galaxy's star formation history (SFH), Hercules is old ($>$ 12 Gyr, with negligible star formation in the last 12 Gyr) and metal-poor ([Fe/H]$ \sim - 2.0$ dex), with an intrinsic spread in metallicity and both [Fe/H] = $-2.3$ dex and $-1.7$ dex populations contributing to the SFH. In the literature there are several spectroscopic and photometric determinations of the metallicity of Hercules. These studies confirm that Hercules shows a spread in metal abundance, with values of the mean metallicity $\langle {\rm [Fe/H]} \rangle$ ranging from about $-$2.0 to $-$2.7 dex \citep{SG07,Kirby08,Koch08,Sand09,Aden09,Aden11}. The line of sight towards the Hercules dSph galaxy is heavily contaminated by Galactic foreground stars, making it hard to determine membership from the CMD alone. Even when radial velocities are added the selection remains uncertain, because the mean velocity of the Hercules dSph galaxy coincides with the velocity of the thick disk (\citealt{Aden09}). Indeed, in the \citet{Aden09} CMD, based on Str\"{o}mgren photometry obtained with the Wide Field Camera of the INT, the RGB of Hercules is not easily identified due to the halo foreground contamination. These authors used the c1 index in the Str\"{o}mgren system to disentangle the galaxy's RGB and HB stars from the foreground contamination. Hercules appears to be highly elongated. The galaxy's structural parameters were obtained by a number of different authors \citep{Bel07,Col07,Martin08,Sand09} who, fitting different stellar profiles to their independent photometric datasets, found similar values for the galaxy's central position, position angle ($\theta$), and ellipticity ($\epsilon$), but rather different half-light radii. In a recent analysis, \citet{Pen09} suggest that the density profiles of relaxed, tidally stripped dwarf spheroidals like Hercules are better approximated by a Plummer law. Moreover, \citet{Sand09} clearly point out the need for deep photometry in order to properly constrain the structural parameters of the new faint MW satellites. In the following, we shall adopt the structural parameters obtained by \citet{Sand09} assuming a Plummer stellar distribution, namely $\theta=-72.59^{\circ}$, $\epsilon =0.67$ and $r_h=6.27\arcmin$. Hercules has no evidence of internal rotation, and a very low velocity dispersion of $\sim 5$ km s$^{-1}$ for \citet{SG07} or 3.72 km s$^{-1}$ for \citet{Aden09}. The explanation of such a large ellipticity in absence of a rotational support might imply that Hercules is not in dynamical equilibrium due to strong tidal distortions \citep[see discussion in][]{Col07,Martin08}. The latest studies estimate distances in the range of $\sim 132$ kpc to $\sim 147$ kpc \citep[see e.g.][]{Col07,Sand09,Aden11} and a total absolute magnitude ranging from $M_V=-6.2\pm 0.4$ mag to $M_V=-6.6\pm0.3$ mag \citep[see e.g.][]{Sand09,Martin08}. The study presented in this paper is based on $B,V$ photometric time-series imaging covering a field of view (FOV) of $\sim 40' \times 36'$ (see Table~\ref{t:obs}), extending well beyond Hercules' half-light radius. These data have allowed us to obtain a complete inventory of the variable stars belonging to the galaxy and to trace the corresponding parent stellar populations. This paper is organized as follows: in Section 2 we present the observations and the data reduction procedures; Section 3 is devoted to the variable stars, whereas the CMDs and the implications for the structure of the Hercules UFD are discussed in Section 4. A new estimate of the distance to Hercules based on the galaxy's RR Lyrae stars is presented in Section 5. Finally, the summary and conclusions in Section 6 close the paper. | In this paper we have presented the first time-series analysis of the Hercules UFD. Using a variety of telescope/instrument facilities we secured $\sim$ 80 and 50 epochs in $B$ and $V$. These data allowed us to detect and characterize 9 RR Lyrae stars (6 ab- and 3 c-type, respectively) and one Anomalous Cepheid. The same observations allowed us to build a deep CMD extending well beyond the galaxy's half-light radius. The main results of this study are listed below: \begin{itemize} \item The average period of the ab-type RR Lyrae stars, $\langle$P$_{ab}\rangle=$0.68 d, qualifies Hercules as an Oosterhoff II system, in good agreement with the vast majority of the UFDs investigated so far. This occurrence favors the hypothesis that the UFDs could be the ``building blocks'' of the Galactic halo, since the pulsation characteristics of their RR Lyrae stars are in agreement with the properties of the MW halo variables. \item Hercules' CMD is dominated by a stellar population at least as old and metal-poor as the GGC M68. This result is in agreement with previous findings. The HB shows some spread. We also detected an overabundance of stars above the HB, thus confirming the previous finding by \citet{Aden09}. This, along with the detection of an Anomalous Cepheid very likely belonging to Hercules, hints at the possible presence of an intermediate-age population about $\sim 2-3$ Gyrs old in Hercules. \item The spatial distribution of Hercules' stars confirms the elongated shape of this galaxy. The signature that Hercules is undergoing tidal disruption is provided by the absence of a clearcut difference between galaxy and field star properties, and by the presence of two RR Lyrae stars lying well beyond the galaxy's half-light radius. \item The RR Lyrae variables were used to obtain independent estimates of the metallicity, reddening and distance to Hercules, for which we find [Fe/H]=$-2.30\pm0.15$ dex, $E(B-V)=0.09\pm0.02$ mag, and $(m-M)_0=20.6\pm0.1$ mag respectively, in very good agreement with literature values. \end{itemize} \bigskip | 12 | 6 | 1206.4031 |
1206 | 1206.3611_arXiv.txt | {Quasars with redshifts greater than 4 are rare, and can be used to probe the structure and evolution of the early universe. Here we report the discovery of six new quasars with $i$-band magnitudes brighter than 19.5 and redshifts between 2.4 and 4.6 from the YFOSC spectroscopy of the Lijiang 2.4m telescope in February, 2012. These quasars are in the list of $z>3.6$ quasar candidates selected by using our proposed $J-K/i-Y$ criterion and the photometric redshift estimations from the SDSS optical and UKIDSS near-IR photometric data. Nine candidates were observed by YFOSC, and five among six new quasars were identified as $z>3.6$ quasars. One of the other three objects was identified as a star and the other two were unidentified due to the lower signal-to-noise ratio of their spectra. This is the first time that $z>4$ quasars have been discovered using a telescope in China. Thanks to the Chinese Telescope Access Program (TAP), the redshift of 4.6 for one of these quasars was confirmed by the Multiple Mirror Telescope (MMT) Red Channel spectroscopy. The continuum and emission line properties of these six quasars, as well as their central black hole masses and Eddington ratios, were obtained. | % \label{sect:intro} The number of known quasars has increased steadily in the past four decades since their discovery in 1963 (Schmidt 1963). In particular, a large number of quasars have been discovered in two large spectroscopic surveys, namely, the Two-degree Field (2dF) survey (Boyle et al. 2000) and the Sloan Digital Sky Survey (SDSS) (York et al. 2000). 2dF mainly selected low redshift ($z<2.2$) quasar candidates with UV-excess (Smith et al. 2005) and has discovered more than 20,000 quasars (Croom et al. 2004). SDSS adopted a multi-band optical color selection method for quasars mainly by excluding the point sources in the stellar locus of the color-color diagrams (Richards et al. 2002) and has identified more than 120,000 quasars (Schneider et al. 2010). 90\% of SDSS quasars have low redshifts ($z<2.2$), though some dedicated methods were also proposed for finding high redshift quasars ($z>3.5$) (Fan et al. 2001a,b; Richards et al. 2002). High-redshift quasars are rare, and those with redshifts greater than 4 represent only 1\% in the total quasar population. In the SDSS DR7 quasar catalog (Schneider et al. 2010), only 1248 (392) among 105783 quasars have redshifts greater than 4 (4.5). Since these $z\sim4$ quasars exist when the universe is at age of 1.57 Gyr, they can be used to probe the structure and evolution of the early universe (Smith et al. 1994; Constantin et al. 2002). In particular, the absorption line spectra of these quasars can give valuable information on the nature of intergalactic medium at high redshift. However, discovering $z\sim4$ quasars is a big challenge because they are fainter than the low redshift quasars due to their larger distances. Moreover, the Ly$\alpha$ emission lines for $z\sim4$ quasars move to the red end of optical spectra, making them hard to be distinguishable from stars due to similar optical colors. Recently, Wu \& Jia (2010) proposed using the $Y-K/g-z$ criterion to select $z<4$ quasars and using the $J-K/i-Y$ criterion to select $z<5$ quasars with the SDSS optical and UKIDSS (UKIRT Infrared Deep Sky Survey)\footnote{The UKIDSS project is defined in Lawrence et al. (2007). UKIDSS uses the UKIRT Wide Field Camera (WFCAM; Casali et al. 2007) and a photometric system described in Hewett et al. (2006). The pipeline processing and science archive are described in Hambly et al. (2008). } near-IR data based on a K-band excess technique (Warren et al. 2000; Hewett et al. 2006; Chiu et al. 2007; Maddox et al. 2008). With these two criteria, we expect to obtain more complete quasar samples than previous ones. Recent optical spectroscopic observations made by the GuoShouJing Telescope (LAMOST) and MMT have demonstrated the success of finding the missing quasars with redshifts between 2.2 and 3 using the Y-K/g-z criterion (Wu et al. 2010a,b; Wu et al. 2011). We also hope to discover some $z\sim4$ quasars with the J-K/i-Y criterion, which is expect to be applicable for selecting the candidates of quasars with redshifts up to 5 (Wu \& Jia 2010). In this letter, we report our discovery of six new high redshift quasars from the spectroscopic observations with the Lijiang 2.4m telescope and MMT in February, 2012. The successful identifications of these high redshft quasars further demonstrate the effectiveness of using our newly proposed criteria for discovering the missing quasars including high-redshift ones. \begin{figure} \centering \includegraphics[width=14.0cm, angle=0]{ms1162fig1.ps} \caption{The YFOSC spectra of six new quasars. From the left to right, the red dashed lines mark the wavelengths of Ly$\alpha$, SiIV and CIV emission lines at the estimated redshift for five $z>3.6$ quasars, while for SDSS J113816.85+045023.6 they mark the wavelengths of CIV and CIII]. } \label{Fig1} \end{figure} | \label{sect:discussion} A complete quasar sample is crucial for studying the large scale structure of the universe. The current available quasar samples are mostly biased towards low redshifts ($z<2.2$) and more efforts are needed to find quasars at high redshift. Wu \& Jia (2010) proposed to obtain a large complete quasar sample with redshifts up to five by combining the $J-K/i-Y$ criterion with the $Y-K/g-z$ criterion to select quasar candidates. Some recent optical spectroscopic observations have demonstrated the success of finding the missing quasars with redshifts between 2.2 and 3 using the $Y-K/g-z$ criterion (Wu et al. 2010a,b; Wu et al. 2011). Our discovery of six high redshift quasars (five with $z>3.6$) from the spectroscopic observations with the Lijiang 2.4m telescope and MMT further demonstrates the effectiveness of using the $ J-K/i-Y$ criterion for discovering quasars with redshifts up to five. Moreover, the identification of five quasars with $z>3.6$ from nine candidates with photometric redshift larger than 3.6 also confirms the robustness of the photometric redshifts estimated by the SDSS and UKIDSS photometric data. We noticed that two among our five $z>3.6$ new quasars do not meet the SDSS $gri$ or $riz$ selection critrion for $z>3.6$ quasars (Fan et al. 2001a,b; Richards et al. 2002), which suggests that about 40\% of such quasars may be missed in the SDSS spectroscopic survey. This obviously needs to be confirmed by future observations of a large sample of $z>3.6$ quasars. Our identification of a $z=4.6$ quasar demonstrates that $z>4$ quasars can be identified with the 2-meter size telescopes in China. We hope more high redshift quasars will be discovered by the future LAMOST quasar survey (Wu 2011), which is aiming at discovering 0.3-0.4 million quasars from 1 million quasar candidates with $i<20.5$, by taking the advantages of 4000 fibers and 5 degree field of view of LAMOST (Su et al. 1998; Zhao et al. 2012). The new quasar selection criteria, such as those based on SDSS, UKIDSS and the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) data (Wu et al. 2012), will be applied for selecting quasar candidates in the LAMOST quasar survey. This will hopefully provide the largest quasar sample in the next five years for further studies of AGN physics, large scale structure and cosmology. \normalem | 12 | 6 | 1206.3611 |
1206 | 1206.6344_arXiv.txt | We study plasma flows above pulsar polar caps using time-dependent simulations of plasma particles in the self-consistent electric field. The flow behavior is controlled by the dimensionless parameter $\alpha=j/c\rhoGJ$ where $j$ is the electric current density and $\rhoGJ$ is the Goldreich-Julian charge density. The region of the polar cap where $0<\alpha<1$ is a ``dead zone'' --- in this zone particle acceleration is inefficient and pair creation is not expected even for young, rapidly rotating pulsars. Pulsars with polar caps near the rotation axis are predicted to have a hollow-cone structure of radio emission, as the dead zone occupies the central part of the polar cap. Our results apply to charge-separated flows of electrons ($j<0$) or ions ($j>0$). In the latter case, we consider the possibility of a mixed flow consisting of different ion species, and observe the development of two-stream instability. The dead zone at the polar cap is essential for the development of an outer gap near the null surface $\rhoGJ=0$. | Magnetic field lines that pass through the light cylinder of a rotating neutron star are twisted and carry electric currents $\bj_B=(c/4\pi)\nabla\times\bB$. These currents are sustained by electric field $E_\parallel$ induced along the magnetic field $\bB$, and ohmic dissipation $E_\parallel j$ feeds the observed pulsar activity. Voltage associated with $E_\parallel$ controls the energies of accelerated particles, creation of secondary electron-positron pairs, and emission of radio waves. The accelerating voltage has been discussed in a number of works on pulsars beginning from early papers in the 1970s (Sturrock 1971; Ruderman \& Sutherland 1975; Arons \& Scharlemann 1979). The key dimensionless parameter of the polar-cap accelerator is \beq \label{eq:alpha} \alpha=\frac{j_B}{c\rhoGJ}, \eeq where $\rhoGJ=-\bOm\cdot\bB/2\pi c$ is the local corotation charge density of the magnetosphere (Goldreich \& Julian 1969). For a special value of $\alpha=\alpha_0$ (close to unity) a steady state was found for the polar-cap flow with significant particle acceleration (e.g. Arons \& Scharlemann 1979; Muslimov \& Tsygan 1992). However, $\alpha$ is not, in general, expected to take this special value (e.g. Kennel et al. 1979). Global solutions for approximately force-free pulsar magnetospheres give $\alpha$ that significantly varies across the polar cap (Timokhin 2006). In general, $\alpha$ can take any value from $-\infty$ to $+\infty$, depending on the polar cap distance from the rotation axis and the location inside the polar-cap region. The character of the polar-cap accelerator strongly depends on $\alpha$ (Mestel et al. 1985; Beloborodov 2008, hereafter B08). The steady solution with $\alpha=\alpha_0\approx 1$ is a separatrix between two opposite regimes of efficient and inefficient acceleration.\footnote{ Hereafter we will refer to this separatrix as $\alpha=1$, neglecting the deviation of $\alpha_0$ from unity. The precise $\alpha_0$ is controlled by the curvature of magnetic field lines and the general relativistic effects (Muslimov \& Tsygan 1992); its exact value is close to unity and is not essential for the rest of the paper.} In particular, if $0<\alpha<1$, $E_\parallel$ is quickly screened in the charge-separated plasma flowing from the polar-cap surface. The electric field satisfies Maxwell equations that read (in the co-rotating frame of the star, see e.g. Fawley et al. 1977; Levinson et al. 2005), \beq \label{eq:Gauss} \nabla\cdot\bE=4\pi(\rho-\rhoGJ), \eeq \beq \label{eq:Maxw} \frac{\partial\bE}{\partial t}=4\pi(\bj_B-\bj). \eeq If $0<\alpha<1$, there exists a velocity $v=\alpha c$ that allows the charge-separated flow $j=\rho v$ to simultaneously satisfy $\rho=\rhoGJ$ and $j=j_B$. If the flow started from the conducting boundary (which has $E=0$) with $v=\alpha c$, no electric field would be generated (then $\nabla\cdot\bE=0$ and $\partial\bE/\partial t=0$). The actual boundary has $v\neq \alpha c$, as charges are lifted from the polar-cap surface with a small initial $v$, comparable to the thermal velocity in the surface material. The deviation of $v$ from $\alpha c$ implies $\rho\neq\rhoGJ$ or $j\neq j_B$, which generates electric field. B08 argued that \Eqs~(\ref{eq:Gauss}) and (\ref{eq:Maxw}) with $0<\alpha<1$ always drive the flow toward $v=\alpha c$, like a pendulum is driven by gravity toward its equilibrium position. The resulting oscillations occur in space or time, according to \Eqs~(\ref{eq:Gauss}) or (\ref{eq:Maxw}), respectively. For example, the steady-state solution for a cold flow exhibits oscillations in space (Mestel et al. 1985; B08). The oscillatory behavior of the flow with $0<\alpha<1$ is, in essence, Langmuir oscillations; they are generated near the boundary where the flow is injected with $v<\alpha c$ and accelerated toward $v=\alpha c$. In this paper, we investigate the accelerator with $0<\alpha<1$ in more detail. In \Sect~2, we write down the steady-state solution for the charge-separated flow, generalized to non-zero temperature of the polar-cap. We argue that the flow is unstable to small perturbations and can develop into a complicated time-dependent state with a broad momentum distribution. To explore the behavior of the flow, we perform fully kinetic time-dependent simulations. The method of simulations is described in \Sect~3, and the results are presented in Sections~4 and 5. Our simulations confirm the predicted turbulent Langmuir oscillations with a small voltage. Particle acceleration in the flow with $0<\alpha<1$ is insufficient to ignite pair creation. Implications of this ``dead zone'' for radio emission and outer gaps in pulsars are discussed in \Sect~6. | We have presented detailed one-dimensional time-dependent simulations of the plasma flow extracted from the polar caps of neutron stars. The simulations provide a fully kinetic description of the flow, with self-consistent electric field and particle distribution function. In this paper, we focused on the regime $0<\alpha<1$, where $\alpha$ is the main parameter of the flow defined by \Eq(\ref{eq:alpha}). In agreement with the estimates of B08, we find that the particles are accelerated to Lorentz factors, \beq \label{eq:gam2} \gamma\approx \frac{1+\alpha^2}{1-\alpha^2}, \eeq and are not capable of igniting pair creation. In this sense, flows with $0<\alpha<1$ are ``dead.'' They are sustained by a modest voltage, oscillating in space and time. Although the simulation is limited to regions close to the pulsar surface, the result does not depend on the simulation box size, and hence should describe the entire polar cap flow, as long as $\alpha$ remains between $0$ and $1$. The parameter $\alpha$ is expected to vary along the magnetic field lines, on a scale comparable to the stellar radius; we have verified that this variation does not change the oscillating behavior of the flow (see also B08). The simulations show how a kinetic instability develops and disrupts the ideal periodic structure found in the analytical models of the dead zone; the instability mechanism is described in \Sects~2 and 4. We find that the momentum distribution function has two distinct parts --- a variable ``cold stream'' and a broad wing at low momenta, which includes particles flowing backward to the polar cap. The fraction of particles in the cold stream is approximately equal to $\alpha$; the remaining fraction $1-\alpha$ forms the broad component. Even though the flow is turbulent, it shows no signs of particle acceleration to energies higher than that of the cold stream. The value of parameter $\alpha$ depends on the location and geometry of the polar cap. A simplest magnetospheric configuration is that of a centered dipole. Then the parameter $\alpha$ depends on the angle between the magnetic and spin axes, $\xi$; besides, it varies across the polar cap. For nearly aligned rotators ($\xi\approx 0$), $0<\alpha<1$ in the central part of the polar cap and $\alpha<0$ in a ring-shaped zone near the edge of the polar cap (Timokhin 2006; Parfrey et al. 2012). In this case, the dead zone occupies the central part of the polar cap, and $e^\pm$ discharge must be confined to the ring, matching the phenomenological ``hollow cone'' model of pulsar emission. In contrast, the polar cap of an orthogonal rotator ($\xi\approx \pi/2$) has $|\alpha|\gg 1$, which enables $e^\pm$ discharge for the entire polar cap. At arbitrary misalignment $0<\xi<\pi/2$, the values of $\alpha$ are provided by global three-dimensional simulations of the magnetospheric structure (e.g. Spitkovsky 2006) and should play a key role for the geometry of the radio beam. We presented our results using plasma skin depth $\lambda_p$ as a unit of length and particle rest-mass $mc^2$ as a unit of energy. In this form, the results do not depend on the charge or mass of the particles extracted from the polar cap, as long as the flow is made of identical particles. In particular, \Eq~(\ref{eq:gam2}) is valid for both electron flow ($\rhoGJ<0$) and ion flow ($\rhoGJ>0$), and the phase-space distribution shown in Figure~5 describes both cases. Note that the accelerating voltage is proportional to the particle mass; voltage implied by \Eq~(\ref{eq:gam2}) is different for ions and electrons by the factor of $m_i/m_e\sim 2\times 10^3$. The relatively high voltage in the ion flow, $e\Phi\approx m_ic^2 (1+\alpha^2)/(1-\alpha^2)$ is still hardly sufficient to ignite $e^\pm$ pair discharge by a seed electron or positron. The identical-particle model may not hold for an ion flow; in this case, new effects may enter the problem. Firstly, heavy ions pulled out from the polar cap may not be completely ionized and begin to lose electrons as they are accelerated and interact with the X-rays above the stellar surface; this process effectively creates new charges, reminiscent of pair creation (e.g. Jones 2012). Secondly, the ion flow may be a mixture of different nuclei which will be accelerated to different Lorentz factors. The mixed ion flow is prone to two-stream instability, possibly leading to formation of plasma clumps and generation of coherent radio emission. In our simulations, we observe the expected two-stream instability, however do not observe significant structure (clumps) in the turbulent flow. This may change in three-dimensional simulations. The frequency of excited waves (comparable to the ion plasma frequency) is in the radio band, and coherent emission from clumps could create bright coherent emission. It remains to be seen whether this mechanism can contribute to the pulsar emission. If it does, it would create an additional component of the radio pulse. In the case of approximately aligned rotator, the additional component would be generated in the central region of the polar cap, leading to a ``hollow cone + core'' structure of the radio pulse. The charge-separated model of the dead zone can be modified to include possible backflowing particles from distant parts of the open field-line bundle (e.g. from a pair-producing outer gap). These particles can contribute to the current density and also serve as an additional background charge density, which may be modeled as a contribution to the effective ``vacuum'' charge density $-\rhoGJ$. This would change the effective $\alpha$ (Lyubarsky 1992; B08), most likely reducing it. An outer gap is expected to form in a charge-separated flow near the null surface $\vec{B}\cdot\vec{\Omega}=0$ (Cheng et al. 1986). On a given field line, the outer gap will be screened if it is loaded by multiple $e^\pm$ pairs produced by the discharge at the polar cap. Thus, the suppression of $e^\pm$ discharge near the field-line footpoint is an essential condition for the existence of an outer-gap accelerator. Therefore, one can expect an outer gap to form on field lines with footpoints in the dead zone. We did not simulate in this paper flows with $\alpha>1$ or $\alpha<0$; in these cases particles must be strongly accelerated. This regime leads to $e^\pm$ discharge that must be unsteady, with a significant intermittent backflow (B08). A model for oscillating discharge may be constructed in hydrodynamical approximation (Levinson et al. 2005), however a fully kinetic description is essential, as demonstrated by our results for the dead zone, where a significant fraction of particles are trapped and form a broad wing at low momenta in the distribution function. The discharge simulation can be done using our setup of a fixed current $j=j_B$ at the stellar surface (BT07) and incorporating pair creation. We defer the simulations with $\alpha>1$ and $\alpha<0$ to a future work. When this work was completed, the preprint by Timokhin \& Arons (2012) came out. They present simulations of charge separated flows, using a similar method, and the results agree with our results for $0<\alpha<1$. They also consider flows with $\alpha<0$ and $\alpha>1$, and find a strong unsteady $e^\pm$ discharge confirming the analysis in B08. | 12 | 6 | 1206.6344 |
1206 | 1206.5751_arXiv.txt | We present deep {\em Giant Metrewave Radio Telescope} (\gmrt) radio observations at 240, 330 and 610 MHz of the complex radio source at the center of the NGC1407 galaxy group. Previous \gmrt\ observations at 240 MHz revealed faint, diffuse emission enclosing the central twin-jet radio galaxy. This has been interpreted as an indication of two possible radio outbursts occurring at different times. Both the inner double and diffuse component are detected in the new \gmrt\ images at high levels of significance. Combining the \gmrt\ observations with archival {\em Very Large Array} data at 1.4 and 4.9 GHz, we derive the total spectrum of both components. The inner double has a spectral index $\alpha=0.7$, typical for active, extended radio galaxies, whereas the spectrum of the large-scale emission is very steep, with $\alpha=1.8$ between 240 MHz and 1.4 GHz. The radiative age of the large-scale component is very long, $\sim 300$ Myr, compared to $\sim$30 Myr estimated for the central double, confirming that the diffuse component was generated during a former cycle of activity of the central galaxy. The current activity have so far released an energy which is nearly one order of magnitude lower than that associated with the former outburst. The group X-ray emission in the \chandra\ and \xmm-Newton images and extended radio emission show a similar swept-back morphology. We speculate that the two structures are both affected by the motion of the group core, perhaps due to the core sloshing in response to a recent encounter with the nearby elliptical galaxy NGC\,1400. | \label{sec:intro} Quoting \cite{1999A&A...348..699L}, {\em ``radio sources are born, grow and finally ... sleep''}. The dormant phase in the evolutionary course of a radio galaxy and fate of the radio-emitting plasma after the cessation of the nuclear activity are among the most intriguing issues in extragalactic astronomy. The active stage of a powerful radio source, typically associated with an elliptical galaxy, can last several $10^7$ up to few $10^8$ years. During this time, the radio galaxy is likely fed by mass accretion onto the supermassive black hole of the host galaxy. Once this accretion is interrupted, or becomes insufficient to support radio activity, the radio source enters a dying phase \citep[e.g.,][and references therein]{2011A&A...526A.148M}; the radio emission passively evolves and rapidly fades, even if expansion losses are negligible and the relativistic electrons are subject only to radiative losses. This is reflected into a pronounced steepening of the integrated radio spectrum whose slope $\alpha$ can reach ultra-steep values, $\alpha \gtrsim 2$ \citep[e.g.,][]{1994A&A...285...27K}, adopting the convention $S_{\nu} \propto \nu^{-\alpha}$ for the synchrotron spectrum, where $S_{\nu}$ is the flux density at the frequency $\nu$. Eventually, the radio emission disappears below the detection limit of present radio telescopes. During the fading stage of a radio source, the central nucleus may switch on again and produce new radio activity, thus leading to a restarted radio source. Evidence for episodic radio activities are reported in the literature for a growing number of radio galaxies (see, for instance, Saikia \& Jamrozy 2009 for a review). Double-double radio galaxies \citep{2000MNRAS.315..371S} qualify as one of the most unmistakable examples of recurrent nuclear activity; here a new pair of inner lobes are produced close to the nucleus before the previously generated, more distant ones have completely faded \citep[e.g.,][]{1996MNRAS.279..257S, 1999A&A...348..699L, 2009BASI...37...63S}. Another type of restarted radio sources are ``nesting'' radio galaxies \citep{2007MNRAS.378..581J}, i.e., sources characterized by a faint, extended structure, within which a bright radio source, extended on a significantly smaller scale, is embedded. Examples are Hercules A \citep{2003MNRAS.342..399G,2005MNRAS.358.1061G}, 3C\,310 \citep{1984ApJ...282L..55V,1986MNRAS.222..753L}, and 4C\,29.30 \citep{2007MNRAS.378..581J}. In few well-studied cases, the large--scale emission and inner source appear clearly, and quite sharply, separated in the spectral index distribution \citep[e.g.,][]{2003MNRAS.342..399G}, with much steeper spectral index values measured for the outer emission. This has been interpreted as evidence for different epochs of jet activity, with the more extended and steeper component being related to an earlier outburst of the central active galactic nucleus (AGN). In this paper, we focus on the complex radio source associated with the elliptical galaxy NGC\,1407 (see Table \ref{tab:ngc1407}), at the center of the poor group of galaxies Eridanus~A at z=0.0059, with a one-dimensional velocity dispersion $\sigma_{v} =$372 km~s$^{-1}$ \citep{2006MNRAS.369.1351B} and 0.3-2 keV X-ray luminosity of $10^{41.7}$ erg~s$^{-1}$ \citep{2006PASA...23...38F}. Based on {\em Giant Metrewave Radio Telescope} (\gmrt) observations at 240 MHz and 610 MHz, Giacintucci et al. (2011, hereinafter G11) suggested that this source may be a restarted radio galaxy. Here, we present a multifrequency radio study of NGC\,1407 based on new, deep \gmrt\ observations at 240 MHz, 330 MHz and 610 MHz, and multifrequency data from the {\em Very Large Array} (\vla) archive. Our study is complemented with the analysis of {\em Chandra} and {\em XMM-Newton} X-ray data. | \label{sec:discussion} \subsection{Restarted activity in NGC\,1407}\label{sec:age} The radio images presented in Section \ref{sec:images} show faint, diffuse emission, approximately 80 kpc across, which wholly encloses and dwarfs a small-scale ($\sim8$ kpc) double source at the center of NGC\,1407. Based on former \gmrt\ observations, G11 suggested that the diffuse outer emission was produced during an earlier cycle of activity of NGC\,1407. The spectral analysis presented in this paper corroborates the multiple outburst scenario. The large-scale emission is found to have an ultra-steep radio spectrum, with $\alpha=1.82$, as typically observed for highly evolved and fading radio sources \citep[e.g.,][]{ 1994A&A...285...27K, 2007A&A...470..875P, 2007A&A...476...99G, 2011A&A...526A.148M}. The spectral age analysis suggests that the radio plasma in such component is at least $\sim$300 Myr old. Furthermore, the radiative model that better describes the spectrum of this emission requires a switch-off of the nuclear engine $\sim$170-190 Myr ago, followed by a dying phase. The inner double has instead a spectral index of $\alpha=0.7$ and a radiative age of $\sim 30$ Myr, consistently with being a currently active and relatively young radio source. The multi-scale radio morphology, combined with the distinct spectral properties of its components, makes NGC\,1407 another remarkable example of a {\em nesting} radio galaxy, other nearby examples being, for instance, 4C29.30 \citep{2007MNRAS.378..581J} and Hercules A \citep{2005MNRAS.358.1061G}. As double-double radio galaxies, these peculiar radio sources provide evidence for recurrent radio activity in elliptical galaxies. In the specific case of NGC\,1407, the small, young double source is currently fed by the central AGN, while the diffuse, steep-spectrum component is associated with relic plasma, which was injected during an earlier radio outburst of the AGN occurred at least $\sim$300 Myr ago. It is important to note that our radiative ages have been estimated neglecting adiabatic expansion. This could be a reasonable approximation for a relaxed, aged plasma in the dying phase such as that in the large-scale component. However, expansion losses can be important for the inner double, as well as for the earlier, active phase of the large-scale component. Indeed, neglecting their effect may lead to an underestimate of the true source age. Our age derivation is also based on the assumption of uniform and constant magnetic field across the source. A calculation of the effects of magnetic field evolution is beyond the purpose of the present paper, and we refer to \cite{1994ApJS...90..955R}, \cite{1999ApJ...512..105J}, \cite{2000AJ....119.1111B}, and references therein, for detailed works on the implications of magnetic field evolution for the source aging. In the case of very aged emission as that in NGC\,1407, the relativistic plasma may be partially mixed with the hot thermal gas. In this case, magnetized filaments can be produced in the radio volume due to plasma instabilities and the total synchrotron emission would result from the convolution of different spectra produced by relativistic electrons emitting in regions with different magnetic field strength \citep[e.g.,][]{2004ApJ...601..778T}. The resulting convolution of the synchrotron kernel with magnetic field intensity and geometry yields a total spectrum which is stretched and thus not straightforwardly related to the spectrum of the emitting electrons \citep[e.g.,][]{1996ApJ...457..150E,1997ApJ...488..146K}. In this case, standard aging analyses, based on the position of the break frequency in the synchrotron spectrum, can give incorrect estimates of the age of the radio emitting electrons. \subsection{Energy output of the radio outbursts} We compared the energy output associated with synchrotron radiation from the two radio outbursts in NGC\,1407. Using the radio luminosity at 240 MHz and $\alpha_{\rm inj}$ in Table 5, we calculated the total radio luminosity over the frequency range 10 MHz--100 GHz adopting the expression in \cite{1987ApJ...316...95O}. We found a bolometric radiative power of $L_{\rm tot} \sim 1.6 \times 10^{42}$ erg s$^{-1}$ for the inner double and $L_{\rm tot} \sim 3.2 \times 10^{42}$ erg s$^{-1}$ for the large-scale emission. Based on our estimates of the radiative age of the two components (Table 5), we found that the current outburst has released an energy of $E_{\rm tot, \, syn} \sim 1.5 \times 10^{57}$ erg so far, while the energy of the former episode of activity, $E_{\rm tot, \, syn} \sim 2.6-3.2\times 10^{58}$ erg, is one order of magnitude larger. It is well known that the synchrotron luminosity is an underestimate of the total energy output of a radio source, which is believed to be dominated by the mechanical work done by the radio jets on the external medium. The mechanical jet power of radio sources can be inferred, for instance, in those groups/clusters with depressions (cavities) in their X-ray surface brightness, interpreted as rising bubbles of relativistic plasma inflated by the central AGN \citep[e.g.,][]{2004ApJ...607..800B,2008ApJ...686..859B,2010ApJ...720.1066C,2011MNRAS.416.2916O}. Here, it is estimated that the ratio of the mechanical to synchrotron luminosities is a few to few thousands for powerful radio sources, and up to several thousand for weaker sources \citep[e.g.,][]{2004ApJ...607..800B}. In those special systems with multiple cavities, interpreted as signature of repeated AGN outbursts, it is also possible to compare the energy outputs of the different episodes of activity. In some systems, such as Hydra A \citep{2007ApJ...659.1153W} and NGC\,5813 \citep{2011ApJ...726...86R}, it is found that the past epoch of activity is the most energetic, suggesting that the mean jet power changed significantly over time or that the current outburst is still ongoing. In other systems (e.g., A\,262, \cite{2009ApJ...697.1481C}), there is evidence for an opposite outburst trend, which may reflect an increase of the fueling of the AGN with time. The current X-ray images of NGC\,1407 do not show evidence of X-ray cavities associated with the large-scale diffuse radio structure. \cite{2010ApJ...712..883D} report a possible small (0.7$\times$0.4~kpc) cavity in the group core, but its identification is dependent on the image processing employed and it appears to be uncorrelated with the small-scale active radio source. A direct measurement of the mechanical power of the radio jets in either period of activity is therefore not possible. If the mean jet mechanical power to radio power ratio has remained constant over the outburst history, NGC\,1407 would then be another system in which the total energy output from the AGN is decreasing with time. On the other hand, the most recent episode of activity may be still ongoing (its radiative age is $\sim$30 Myr), and it is possible that the jet power may increase. \subsection{Pressure comparison} Assuming energy equipartition arguments, we can derive the non-thermal pressure in the radio source and compare it to the pressure of the surrounding X-ray gas. Under the assumptions listed at the beginning of Sect.~\ref{sec:spectra} (see point 3), and adopting $\alpha_{\rm inj}$ and $B_{\rm eq}$ in Table 5, we calculated a radio pressure in the large-scale component of $P_{\rm radio, diffuse}= 2.5 \times 10^{-13}$ erg cm$^{-3}$ and $P_{\rm radio, S3}= 2.8 \times 10^{-12}$ erg cm$^{-3}$ for the inner double. Even though the X-ray pressure profile in Fig.~\ref{fig:press} does not cover the whole extent of the radio emission, it is clear that the pressure of the X-ray gas is at least one order of magnitude higher than the radio pressure, with values ranging from $\sim 10^{-11}$ erg cm$^{-3}$ at $\sim$10 kpc from the center to $\sim5 \times 10^{-12}$ erg cm$^{-3}$ at $\sim 40$ kpc. This is not unusual for cool-core systems, where a similar pressure discrepancy is often found \citep[e.g.,][]{2005MNRAS.364.1343D}, suggesting a departure from equipartition conditions or an additional pressure support in the radio lobes, for instance, in the form of thermal gas \citep[e.g.,][]{2010MNRAS.407..321O}. Alternatively, an energetically dominant population of non-radiating relativistic particles (protons) may contribute to the internal pressure. The achievement of pressure balance in NGC\,1407 would then require a ratio of energy in non-radiating particles to the energy in electrons $k \sim 100-1000$. These values are in the range typically found for samples of centrally-located radio galaxies in groups and clusters \citep{2004MNRAS.355..862D,2008ApJ...686..859B,2008MNRAS.386.1709C,2010ApJ...714..758G} and suggest that that the lobes have been fed by {\em heavy} jets or have entrained thermal material during their propagation through the cluster/group atmosphere \citep[e.g.,][]{2010MNRAS.407..321O}. The diffuse and distorted morphology of the large-scale emission in NGC\,1407, combined with the very long age inferred for this component, suggest that a mixing of the radio plasma and ambient thermal gas has already occurred at some level, favoring the entrainment scenario. \subsection{Motion of the galaxy} The morphology of NGC\,1407 in the X-ray band (Fig.~\ref{fig:xmm} and Fig.\ref{fig:chandra}) suggests that the galaxy is in motion. A surface brightness edge is visible in the north side of the galaxy with gas swept back to the east and west into two cool wings. There is some indication of a possible discontinuity in the surface brightness profile of the northern quadrant of the galaxy at this position (Fig.~10). If confirmed, this feature would suggest the presence of a shock or a cold front. The Fe-peak temperature map shows variations across the edge which seem consistent with the the hypothesis of a cold front, i.e., with the cooler gas behind the front (Fig.~12). However, the available data are not deep enough to constrain the temperature jump across the edge and thus the nature of this feature. Overall, NGC\,1407 appears to be moving northward with its halo being stripped by the surrounding IGM. The location of the putative front is consistent with this scenario. If the edge is a cold front, it may indicate that the gas is sloshing in response to a recent interaction, such as a close passage, of NGC\,1400. The X-ray trail visible in the \xmm\ image (Fig.~8) could originate from the same interaction. The diffuse radio structure also shows a swept-back shape, but on a larger scale, similar to wide-angle tail (WAT) radio galaxies, although no jets are present. As discussed above, this emission is found to be the remnant of a former radio outburst of the central galaxy, consistent with the absence of jets. The X-ray wings seem to be inside the radio emission, and the radio contours appear compressed at the northern edge (e.g., Fig.\ref{fig:temp}). The similarities in the radio and X-ray morphologies suggest that the two structures are co-spatial and both affected by the motion of the galaxy. The X-ray wings may then represent a wake of cooler galactic material which has been stripped by the same ram-pressure that bent the diffuse radio emission. There are several ways in which we can obtain rough estimates of the velocity of NGC\,1407 relative to the IGM. If we assume, for simplicity, that the bending of the diffuse radio emission is caused by ram pressure, we can use the Euler equation in the form $$\frac{\rho_{\rm radio} v_{\rm radio}^2}{r_c} = \frac{\rho_{\rm IGM} v_{\rm gal}^2}{r_{\rm radio}}$$ \citep[e.g.,][]{1985ApJ...295...80O}, where $r_c$ is their curvature radius, $r_{\rm radio}$ is the radius of the radio {\em tails}, $\rho_{\rm radio}$ and $v_{\rm radio}$ are the density and velocity of the radio-emitting plasma, $v_{\rm gal}$ is the velocity of the galaxy relative to the IGM, and $\rho_{\rm IGM}$ is the density of the IGM. From the images presented in Sect.~\ref{sec:images}, we estimate $r_{\rm radio} \sim 10$ kpc and $r_c \sim 30$ kpc. Based on the spatial extent of the diffuse radio emission (LLS$\sim 80$ kpc) and age of the active phase $t_{\rm CI} \sim 90-130$ Myr (Table 5), we derive a first order estimate of the growth velocity of the radio source of $v_{\rm radio} = {\rm LLS}/t_{\rm CI} \sim 0.002c-0.003c$, where $c$ is the speed of light. Finally, from the \chandra\ data, we obtain $\rho_{\rm IGM} \sim 0.0015$ cm$^{-3}$ within the central 50 kpc and assume $\rho_{\rm radio} \sim 10^{-3} \rho_{\rm IGM}$ \citep[e.g.,][]{2011ApJ...743..199D}. This gives a very low velocity, only $v_{\rm gal} \sim$ 20 km s$^{-1}$. Alternatively, we can make a simple calculation assuming that the galaxy is moving north, that the radio outburst was located where we now see the southern boundary of the diffuse structure, and that it has moved at a constant velocity since that time. Neglecting projection effects and basing our estimate on the 240~MHz map, the galaxy would then have traveled 28-35~kpc over a period of $\sim$300~Myr, suggesting a velocity of 90-115 km s$^{-1}$. These velocity estimates conflict to some extent with the X-ray morphology. If a cold front is present in NGC~1407, we would expect the galaxy to be moving at a significant fraction of the sound speed. For a temperature of 1.1~keV, as is found immediately north of the surface brightness edge, the sound speed is $v_{sonic}\sim$500 km s$^{-1}$, so a velocity of at least 250 km s$^{-1}$ would be likely. However, there are large uncertainties on our estimates. If sloshing is taking place, the most likely cause is perturbation by a tidal encounter with NGC\,1400. Given the 1100 km s$^{-1}$ velocity offset between NGC\,1400 and the group mean, it is clear that sloshing motions would include a significant line of sight component. Since we see signs of a surface brightness edge in the X-ray and compressed contours in the radio, the velocity in the plane of the sky is greater than or comparable to that in the line of sight, but this still leaves an uncertainty of up to a factor $\sqrt{2}$ in the true velocity of NGC\,1407. The velocity could thus be as high as 160 km s$^{-1}$, or 0.3$v_{sonic}$. Projection effects could have an impact on our estimate of the ram-pressure timescale. If the diffuse structure is in fact two old radio lobes, the axis of the jet which formed them may have been at an angle to the line of sight. In this case the lobes would be at a larger radius than we have assumed, and would experience lower external pressures. Mixing of entrained thermal plasma into the lobes would raise their effective density. Both of these factors could increase the ram--pressure velocity estimate, though it is difficult to place limits on their impact. As discussed above, the X-ray images of NGC\,1407 suggest that the group is not relaxed. A connection between bending of radio jets and sloshing induced by minor mergers has been proposed in cluster cores \citep[e.g.,][]{2004ApJ...616..178C,2006ApJ...650..102A,2012arXiv1203.2312M}, and it seems possible that NGC\,1407 is an example of this process at work in a galaxy group. In particular, we speculate that the group core is sloshing along the North-South axis, in response to a possible recent interaction with the nearby group galaxy NGC\,1400. The gas motions induced by such sloshing may be then shaping both the X-ray and radio structures. There is a degree of tension between the velocities estimated from the radio and X-ray data. Resolving this issue would require deeper X--ray data, capable of detecting any cold front and providing a more reliable velocity measurement, and of determining the filling factor of the diffuse radio structure. | 12 | 6 | 1206.5751 |
1206 | 1206.5798_arXiv.txt | We present EXOFAST, a fast, robust suite of routines written in IDL which is designed to fit exoplanetary transits and radial velocity variations simultaneously or separately, and characterize the parameter uncertainties and covariances with a Differential Evolution Markov Chain Monte Carlo method. We describe how our code incorporates both data sets to simultaneously derive stellar parameters along with the transit and RV parameters, resulting in more self-consistent results on an example fit of the discovery data of HAT-P-3b that is well-mixed in under five minutes on a standard desktop computer. We describe in detail how our code works and outline ways in which the code can be extended to include additional effects or generalized for the characterization of other data sets -- including non-planetary data sets. We discuss the pros and cons of several common ways to parameterize eccentricity, highlight a subtle mistake in the implementation of MCMC that could bias the inferred eccentricity of intrinsically circular orbits to significantly non-zero results, discuss a problem with IDL's built-in random number generator in its application to large MCMC fits, and derive a method to analytically fit the linear and quadratic limb darkening coefficients of a planetary transit. Finally, we explain how we achieved improved accuracy and over a factor of 100 improvement in the execution time of the transit model calculation. Our entire source code, along with an easy-to-use online interface for several basic features of our transit and radial velocity fitting, are available online at http://astroutils.astronomy.ohio-state.edu/exofast. | In the mere 17 years since the first discovery of a ``Hot Jupiter'' around a main sequence star \citep{mayor95}, the study of exoplanets has exploded into one of the most vibrant and rapidly developing fields in astronomy today. Over 500 exoplanets have been confirmed, and four times that many candidates have been identified \citep{batalha12}. The pace of exoplanet discoveries has consistently increased, with new exoplanet search methods continuously being developed and implemented, and new and surprising classes of planets routinely being uncovered as new regimes of parameter space are explored. Meanwhile, an enormous amount of effort is being put into developing theories of planet formation and evolution that can encompass the astonishing diversity of planetary systems that has emerged. The transit method has been at the center of this revolution, not only because it has expanded the region of parameter space to which we are sensitive, but more importantly, it can provide a seemingly endless wealth of information about each planet (see \citealt{winn10a} for a comprehensive review). For example, precise photometry during the primary transit can be used to measure the planet radius and orbital inclination, and when combined with the minimum mass inferred from radial velocity (RV) studies, yield the true planet mass and average density, thereby constraining the planet's structure \citep{guillot05, sato05, charbonneau06, fortney06}. Photometric observations during both primary transits and secondary eclipses enable the study of their atmospheres \citep{charbonneau02, vidal03} and thermal emission \citep{deming05, charbonneau06, deming06}. Variations in the timing and shape of the eclipses and transits hint at the existence of other bodies in the system \citep{miralda02, holman05, agol05, steffen05, ford06b, ford07}, constrain orbital evolution due to tides or other effects \citep{sasselov03, fabrycky08, hellier09}, probe ``weather'' in exoplanet atmospheres \citep{rauscher07}, and provide a probe of the interior structure and oblateness of exoplanets \citep{ragozzine09, carter10}, to name a few. In resonant configurations, even Eris-mass planets may be detectable via such Transit Timing Variations (TTVs) using current ground-based technology for favorable systems \citep{carter11}. Further, the projected angle between the spin axis of the star and the orbit of the planet can be measured via spectroscopic observations during transit to provide diagnostic information of the physical processes at work in the migration of Hot Jupiters \citep{queloz00, winn05, gaudi07, triaud10}. Indeed, the combination of radial velocity and transit data provides the most thorough insights into a planetary system of any demonstrated planet detection and characterization method to date. Because of the wealth of information that can be derived from these planets, it is important to carefully consider the best ways to extract such information from the data sets we acquire such that results are limited by the data and can be compared in a manner that is as homogeneous and consistent as possible. The Markov Chain Monte Carlo method has become a standard tool in exoplanet research \citep[e.g.][]{ford05, gregory05, ford06a, winn10b}, and has recently begun to be replaced with a faster, more elegant flavor: Differential Evolution MCMC, DE-MC \citep[e.g.,][]{braak06,johnson11, carter11b, doyle11}. Many public codes exist to model transit light curves and/or radial velocities, including TAP \citep{gazak11}, JKTEBOP, \citep{southworth08}, FITSH \citep{pal12}, PHEOBE \citep{prsa05}, VARTOOLS \citep{hartman08}, Nightfall\footnote{http://www.hs.uni-hamburg.de/DE/Ins/Per/Wichmann/Nightfall.html}, PhoS-T \citep{mislis12}, and Systemic \citep{meschiari09}. The goal of this paper is to present an additional code, EXOFAST, which we believe provides a valuable combination of features not present in any currently-public code: simultaneous and self-consistent radial velocity, transit, and stellar parameter fitting; fast, robust, DE-MC characterization of errors; intuitive outputs, careful attention to realistic priors; non-interactive (easy to pipeline); well documented; and easy to install, use, and customize. Providing a completely general code that can fit any conceivable planetary phenomenon without modification is not practical. Rather than attempt to be comprehensive, our goal was to provide a modular, easily-extensible framework with a relatively straightforward but powerful example implementation for exoplanets that fits a single-planet system which has either or both RV and primary transit data. This framework and example implementation can be adapted to add additional effects as the data are able to constrain them (e.g., TTVs, secondary eclipses), impose different priors, or even analyze completely different problems (e.g., Supernovae, Cepheids). While IDL is a proprietary language that is generally slower than low-level languages like C and Fortran, we chose to use this language because of the large library of existing code, the ease of development, and the fact that well-written IDL is comparable in speed to higher level languages for most (i.e., non-serial) applications. Of course, an MCMC code is necessarily serial (i.e., one cannot calculate an arbitrary step in the Markov Chain without first calculating the step before it), but the vast majority of the time spent is the model evaluation at each step, which has been carefully vectorized whenever possible. For those unable or unwilling to purchase an IDL license, the GNU Data Language (GDL)\footnote{http://gnudatalanguage.sourceforge.net/} is an open-source compiler that claims full syntax compatibility with code up to IDL version 7.1. Our code does not work out of the box with GDL, but some users have gotten the core features working. Future updates will keep compatibility with GDL in mind. In \S\ref{sec:overview}, we provide a brief summary of the general problem of fitting data sets, an overview of how MCMC works, and why it is preferred over alternative methods of fitting data and estimating uncertainties. The discussion here and routines cited are completely general to the problem of fitting any model to a data set and properly characterizing the uncertainties -- it is not just applicable to exoplanets or even just astronomy. Next, we describe our specific procedure to fit RV data (\S\ref{sec:rv}), including a detailed discussion of different ways to parameterize eccentricity (\S\ref{sec:eparam}), and two potentially-common mistakes when using MCMC, both of which can inflate the measured eccentricity of intrinsically circular orbits significantly. We discuss our procedure to fit transit data in \S\ref{sec:transit}, and combined RV and transit data sets simultaneously in \S\ref{sec:rvtran}. Section~\ref{sec:example} walks through an example fit of HAT-P-3b with real, public data to explain how the code works, what it does, and what its outputs are. Our online interfaces to the most useful codes are presented in \S\ref{sec:online}. Along with these online interfaces, all of the source code described here is available online\footnote{http://astroutils.astronomy.ohio-state.edu/exofast/}. Those already familiar with MCMC and the basics of light curve and RV modeling may find it most efficient to begin with the discussion on eccentricity parameterization (\S\ref{sec:eparam}), then skip to \S\ref{sec:rvtran}, where we discuss our unique approach to fitting RV and Transit data simultaneously. Appendix~\ref{sec:analyticld} demonstrates that the linear and/or quadratic limb darkening coefficients can be fit analytically, which can reduce the dimensionality of a non-linear solver, thereby drastically increasing its speed. Specific improvements to the \citet{mandel02} code to calculate the quadratically limb-darkened flux during transit, including a factor of $\sim$100 improvement in speed that cuts the run time of typical RV and transit fit from an hour to a few minutes, are described in appendix~\ref{sec:occultquad}. Appendix~\ref{sec:random} discusses a problem with IDL's built-in random number generator and provides an alternative at a moderate increase in the overall run time. Appendix~\ref{sec:nege} discusses a way to interpret a negative eccentricity that provides continuous models across the boundary at e=0. Lastly, appendix~\ref{sec:runtime} discusses the execution time and identifies areas for future improvement. | 12 | 6 | 1206.5798 |
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1206 | 1206.0342_arXiv.txt | Hydrocarbons are ubiquitous in the interstellar medium, observed in diverse environments ranging from diffuse to molecular dark clouds and strong photon-dominated regions near HII regions. Recently, two broad diffuse interstellar bands (DIBs) at 4881\AA\ and 5450\AA\ were attributed to the linear version of propynylidene \linear, a species whose more stable cyclic conformer \cyclic\ has been widely observed in the diffuse interstellar medium at radio wavelengths. This attribution has already been criticized on the basis of indirect plausibility arguments because the required column densities are quite large, N(\linear)/\EBV\ $= 4 \times 10^{14}~\pcc$ mag$^{-1}$. Here we present new measurements of N(\linear) based on simultaneous 18-21 GHz VLA absorption profiles of cyclic and linear C$_3$H$_2$taken along sightlines toward extragalactic radiocontinuum background sources with foreground Galactic reddening \EBV\ = 0.1 - 1.6 mag. We find that N(\linear)/N(\cyclic) $\approx 1/15 - 1/40$ and N(\linear)/\EBV\ $\approx 2\pm 1 \times10^{11}\pcc$ mag$^{-1}$, so that the column densities of \linear\ needed to explain the diffuse interstellar bands are some three orders of magnitude higher than what is observed. We also find N(\cfh)/\EBV\ $<1.3 \times10^{13}\pcc$ mag$^{-1}$ and N(\anicfh)/\EBV\ $< 1\times 10^{11}\pcc$ mag$^{-1}$ ($3\sigma$). Using available data for CH and \cch\ we compare the abundances of small hydrocarbons in diffuse and dark clouds as a guide to their ability to contribute as DIB carriers over a wide range of conditions in the interstellar medium. | The cyclic (ring) conformer of propynylidene, \cyclic, was discovered and identified in the interstellar medium (ISM) by \cite{MatIrv85} and \cite{ThaVrt85} and quickly recognized as a ubiquitous tracer of molecular gas in both dark \citep{MadIrv+89,CoxWal+89} and diffuse clouds \citep{CoxGue+88}. It is abundant even in strong photon-dominated regions (PDR) like the Horsehead Nebula \citep{PetTey+05}. Extensive recent surveys of the small hydrocarbons \cyclic\ and \cch\ in mm-wave absorption \citep{LucLis00C2H,GerKaz+11} show that they have a nearly-fixed relative abundance N(\cch)/N(\cyclic) $\approx$ 20 in diffuse gas, much above that seen in the dark cloud TMC-1 where N(\cch)/N(\cyclic) $\approx 2$ \citep{OhiIrv+92}, and a nearly constant abundance with respect to \HH, X(\cyclic) = N(\cyclic)/N(\HH) $\approx 2-3\times 10^{-9}$. By contrast, when the less stable linear conformer \linear\ was detected in the ISM \citep{CerGot+91} it was found to be much less abundant than \cyclic\ in the dark cloud TMC1, with N(\cyclic)/N(\linear) $\approx$ 100 and N(\linear) = $2.5\times10^{12}\pcc$. Subsequent observations of \linear\ in mm-wave absorption from diffuse clouds in front of the distant HII regions W49 and W51 \citep{CerCox+99} found that the ratio \cyclic/\linear\ was smaller, N(\cyclic)/N(\linear) = 3-7, although with small column densities overall, N(\linear) $\la 5\times10^{11}\pcc$. Despite strong suggestions that column densities of \linear\ are modest in the ISM, whether in diffuse or dark gas, \cite{MaiWal+11} recently hypothesized that broad diffuse interstellar bands (DIBs) at 4881\AA\ and 5450\AA\ are produced by a ground-state electronic transition of \linear, with N(\linear) $\approx 2-5\times10^{14}$ toward HD 206267 (\EBV\ = 0.52 mag) and HD 183143 (\EBV\ = 1.27 mag). This assignment has already been criticized for obvious reasons based on indirect plausibility arguments \citep{KreGal+11,OkaMcC11}. Here we show directly that the diffuse ISM is in no way anomalous with regard to the \linear\ abundance. Seen in absorption originating in Galactic diffuse and translucent clouds occulting 4 compact extragalactic mm-wave continuum sources along sightlines with \EBV\ = 0.1 mag - 1.6 mag, column densities of \linear\ are indeed small, with N(\linear) $\la 1 - 3\times10^{11}\pcc$. The plan of this work is as follows. In Section 2 we describe new 18-21 GHz absorption line observations of \linear, \cyclic, \cfh\ and \anicfh. In Section 3 we discuss the abundances and column densities of these and other small hydrocarbons so that their ability to contribute as carriers of DIBs may be accurately assessed. | The sightlines in this work substitute distant radio sources for early-type stars used in optical/uv absorption-line studies, uniquely providing access to a diverse set of polar polyatomic molecules. Nonetheless there are comparable column densities of all species that are accessible in both wavelength regimes, OH, CO and CN \citep{LisLuc96,LisLuc98,LisLuc01} and comparable thermal pressures as judged in the radio by the rotational excitation of CO \citep{LisLuc98}. CO column densities are small and the gas-phase carbon is inferred to reside overwhelmingly in C\p\ even when N(C\p) is not measured directly, implying that the ionization fraction is similar even if the flux of ionizing photons will occasionally be larger in optical absorption line studies. Although there is no theory that reliably predicts the abundances of polyatomic molecules in diffuse clouds, similar circumstances should affect both \linear\ and DIBs no matter which measurement technique is employed The simple hydrocarbons observed and discussed here generally have well-defined abundances with respect to reddening and \HH\ over a wide range of conditions, especially CH, \cch\ and \cyclic, with slightly larger variations for \linear. Bulk effects naturally introduce some degree of superficial correlation with \EBV\ even for unrelated species that are mixed into the ISM along the line of sight, but proportionality between the column densities of small hydrocarbons and N(\HH) reflects a concentration within regions of high molecular fraction, typically \fH2 $\approx$ 0.4 as noted above. This may make the small hydrocarbons rather poorly suited to be candidate carriers for the most heavily-studied DIBs \citep{FriYor+11,VosCox+11}, which do not have obvious associations with N(\HH) and are more usually attributed to regions of low molecular fraction. In any case, the column density per unit reddening of \linear\ in diffuse clouds is a factor 2000 below the value N(\linear)/\EBV\ $= 4 \times 10^{14}~\pcc$ mag$^{-1}$ that was hypothesized by \cite{MaiCha+11} in order that \linear\ could be the carrier of the DIBs at 4881\AA\ and 5450\AA. The most abundant hydrocarbons in diffuse gas, as toward TMC-1, are CH and \cch\ with X(CH) = $4 \times 10^{-8}$ and X(\cch) $= 6\times 10^{-8}$. Taken relative to \cch\ in diffuse gas one has X(\cyclic)/X(\cch) = 1/21 and X(\linear)/X(\cch) = 1/350 from the mean values given in Table 3, with X(\cfh)/X(\cch) $< 1/7$ in the two best cases and X(\anicfh)/X(\cch) $ \la 1/1000$ quite generally. \cite{GreCar+11} review the history of (failed) DIB attributions, including \linear\ and the equally recent example of H-\cfh-H\p, proposed by \cite{KreBel+10} and criticized by \cite{MaiCha+11}. Other attributions recently deemed to be insupportable include heavier species such as the simplest PAH naphthalene (C$_{10}$H$_8$), anthracene (C$_{14}$H$_{10}$) and their cations \citep{GalLee+11}. Perhaps still hanging in the balance, but by no means generally accepted, is a tentative attribution to C\HH CN$^-$\ by \cite{CorSar07}. Species as complicated as naphthalene are likely to prove as elusive at radio wavelengths as in the optical/uv domain, but the sort of work performed here may be straightforwardly extended to other polar species such as C$_3$H and C\HH CN, especially considering the modest observing times required and the increasing availability of wider instantaneous bandwidths at the GBT and VLA. Searches may also be possible for non-polar species that are observable in their magnetic dipole transitions as recently discussed by \cite{MorMai11}. | 12 | 6 | 1206.0342 |
1206 | 1206.0497_arXiv.txt | We describe observations of a white-light flare (SOL2011-02-24T07:35:00, M3.5) close to the limb of the Sun, from which we obtain estimates of the heights of the optical continuum sources and those of the associated hard X-ray sources. For this purpose we use hard X-ray images from the \textit{Reuven Ramaty High Energy Spectroscopic Imager (RHESSI)}, and optical images at 6173~\AA~from the \textit{Solar Dynamics Observatory (SDO)}. We find that the centroids of the impulsive-phase emissions in white light and hard X-rays ($30-80$~keV) match closely in central distance (angular displacement from Sun center), within uncertainties of order 0.2$''$. This directly implies a common source height for these radiations, strengthening the connection between visible flare continuum formation and the accelerated electrons. We also estimate the absolute heights of these emissions, as vertical distances from Sun center. Such a direct estimation has not been done previously, to our knowledge. Using a simultaneous 195~\AA~image from the \textit{Solar-Terrestrial RElations Observatory (STEREO-B)} spacecraft to identify the heliographic coordinates of the flare footpoints, we determine mean heights above the photosphere (as normally defined; $\tau = 1$ at 5000~\AA) of $305 \pm 170$~km and $195 \pm 70$~km, respectively, for the centroids of the hard X-ray (HXR) and white light (WL) footpoint sources of the flare. These heights are unexpectedly low in the atmosphere, and are consistent with the expected locations of $\tau = 1$ for the 6173~\AA~and the $\sim$40~keV photons observed, respectively. | \label{sec:intro} The white-light continuum of a solar flare (WLF) was the first manifestation of a flare ever detected \citep{1859MNRAs..20...13C,1859MNRAs..20...16H}. Nevertheless its origin has remained enigmatic over the intervening centuries. This continuum contains a large fraction of the total luminous energy of a flare \citep[e.g.,][]{1989SoPh..121..261N}, and so its identification has always posed an important problem for solar and stellar physics. Although the most obvious flare effects appear in the chromosphere and corona, the physics of the lower solar atmosphere has great significance for the reasons described by Neidig. The association of white-light continuum with the impulsive phase of a solar flare has long been known, and cited as an indication that high-energy nonthermal particles have penetrated deep into the lower solar atmosphere \citep{1970SoPh...15..176N,1970SoPh...13..471S}. Such an association would link the acceleration of ``solar cosmic rays,'' as they were then known, with the intense energy release of the impulsive phase, and this association has proven to be crucial to our understanding of flare physics \citep{1976SoPh...50..153L,2003ApJ...595L..69L,2005JGRA..11011103E,2007ApJ...656.1187F}. Because the continuum appears in the emission spectrum, and because this nominally originates in an optically-thick source, this energy release may drastically distort the structure of the lower solar atmosphere during a flare. Accordingly the usual methods for modeling the atmosphere \citep[e.g.,][]{1981ApJS...45..635V}, which assume hydrostatic equilibrium, may be too simple. \begin{figure*}[htpb] \includegraphics[width=0.95\textwidth]{fig1.eps} \caption{ HMI intensity continuum difference image (prime, 07:31:13.40~UT; reference, images from 07:25-07:28~UT) combined with white-light difference and RHESSI CLEAN contour plots (red and blue, respectively). The Hard X-ray images (30-80 keV) made with the CLEAN technique for the interval 07:30:50.9 -- 07:31:35.9~UT, exactly that of the HMI cadence. These images were made with \textit{RHESSI} subcollimators 1-4, with uniform weighting, giving an angular resolution (FWHM) of 3.1$''$. The orange contours in panel a) show the 6-8 keV Soft X-ray source at the same time, defining a loop structure connecting the footpoints. The dotted line shows the locus of the STEREO/EUVI source positions, which in this projection show the projected angular location of the photosphere (see Section~\ref{sec:absolute}).} \label{fig:footpoints} \end{figure*} The association with hard X-ray bremsstrahlung has always suggested $>$10~keV electrons in particular \citep{1972SoPh...24..414H,1975SoPh...40..141R,1992PASJ...44L..77H,1993SoPh..143..201N}. According to the standard thick-target model the primary particle acceleration occurs in the corona above the flare sources. The presence of fast electrons automatically implicates the chromosphere as well, rather than the photosphere, because of the relatively short collisional ranges of such particles. Nevertheless the indirect nature of the bremsstrahlung emission mechanism \citep[e.g.,][]{1971SoPh...18..489B} has made it difficult to rule out other processes, such as energy transport by protons \citep{1969sfsr.conf..356E,1970SoPh...15..176N,1970SoPh...13..471S}. The bremsstrahlung signature may also result from acceleration directly in the lower atmosphere, rather than in the corona \citep{2008ApJ...675.1645F}. Finally radiative backwarming could also provide a mechanism for explaining for the observed tight correlation of hard X-rays and white-light continuum \citep{1989SoPh..124..303M}. This mechanism involves irradiation and heating of the photosphere, rather than the chromosphere, as would correspond to the shorter stopping distance of energetic electrons in a thick-target model \citep{1972SoPh...24..414H}. In the modern era we are seeing a rapid increase in our understanding of these processes, thanks to the excellent new data from various spacecraft and ground-based observatories. This paper reports on a first good example of a flare near the limb, observed at high resolution both in hard X-rays by \textit{RHESSI}, and at 6173~\AA~in the visible continuum by the \textit{Solar Dynamics Observatory (SDO)} spacecraft via its Helioseismic Magnetic Imager (HMI) instrument \citep{2012SoPh..275..229S}. This event, SOL2011-02-24T07:35 (M3.5), occurred just inside the limb (NOAA coordinates N14E87), so that source heights could be compared by simple projection. We analyze data from \textit{RHESSI}, HMI, and the Extreme Ultraviolet Imager (EUVI) on \textit{STEREO-B)}. \cite{2011A&A...533L...2B} have already studied this flare, using the same data but without reference to the \textit{STEREO-B} observations. | \label{sec:concl} In this study we have compared hard X-ray and white-light observations of a limb flare, SOL201-02-24T07:35:00. The relative positions of the sources agree well, for each of the double-footpoint sources; since the local vertical maps almost exactly onto the solar radial coordinate, this means that the source heights match well. Our uncertainties on this centroid matching are of order 0.2$''$. This result strongly associates the white-light continuum with the collisional losses of the non-thermal electrons observed via bremsstrahlung hard X-rays in the impulsive phase of the flare. We have also used the EUVI data in the 195~\AA~band from \textit{STEREO-B} to determine the heliographic coordinates of the flare footpoints, from a near-vertical vantage point. To our knowledge, this enables the first direct determination of the absolute height of a white-light flare and its associated hard X-ray sources. Surprisingly, our estimates lie close to (if not below) the projected heights of optical depth unity for both 40~keV hard X-rays and for optical continuum at 6173~\AA. They also lie well below the expected penetration depth of the $\sim$50-keV electrons needed to produce the hard X-rays (about 800~km for the Fontenla et al. quiet-Sun atmosphere, vs. 200-400~km as observed). At present we have no explanation for this striking result and do not speculate about it, because it depends upon only a single flare event. We are sure that there are other comparable events in the existing data and hope to see a generalization of these results based on similar analyses. \bigskip\noindent {\bf Acknowledgements:} This work was supported by NASA under Contract NAS5-98033 for \textit{RHESSI} for authors Hudson, Hurford, Krucker, Lin, and Mart{\' i}nez Oliveros. R. Lin was also supported in part by the WCU grant (R31-10016) funded by the Korean Ministry of Education, Science, and Technology. Jesper Schou and Sebastien Couvidat are supported by NASA contract NAS5-02139 to Stanford University. The HMI data used are courtesy of NASA/SDO and the HMI science team. We thank Martin Fivian for helpful discussions of the RHESSI aspect system. Alex Zehnder's precise metrology of RHESSI has made this analysis possible in the first place. We further thank M. Waltham for comments on the saturation properties of the SECCHI CCD detectors. \nocite{2011A&A...533L...2B} \nocite{1971SoPh...18..489B} \nocite{1859MNRAs..20...13C} \nocite{1969sfsr.conf..356E} \nocite{2005JGRA..11011103E} \nocite{2001SoPh..204...69F} \nocite{1859MNRAs..20...16H} \nocite{1972SoPh...24..414H} \nocite{1992PASJ...44L..77H} \nocite{1976SoPh...50..153L} \nocite{2003ApJ...595L..69L} \nocite{1976ApJ...203..753L} \nocite{1989SoPh..124..303M} \nocite{1970SoPh...15..176N} \nocite{1989SoPh..121..261N} \nocite{1993SoPh..143..201N} \nocite{1975SoPh...40..141R} \nocite{1970SoPh...13..471S} \nocite{1981ApJS...45..635V} \nocite{2009RAA.....9..127W} \nocite{2010ApJ...715..651W} \nocite{1998ApJ...500L.195B} | 12 | 6 | 1206.0497 |
1206 | 1206.0032_arXiv.txt | We present results of optical and infrared photometric monitoring of the eclipsing low-mass X-ray binary V395 Car (2S~0921--630). Our observations reveal a clear, repeating orbital modulation with an amplitude of about one magnitude in $B$, and $V$ and a little less in $J$. Combining our data with archival observations spanning about 20\,years, we derive an updated ephemeris with orbital period $9.0026\pm0.0001$\,d. We attribute the modulation to a combination of the changing aspect of the irradiated face of the companion star and eclipses of the accretion disk around the neutron star. Both appear to be necessary as a secondary eclipse of the companion star is clearly seen. We model the $B$, $V$, and $J$ lightcurves using a simple model of an accretion disk and companion star and find a good fit is possible for binary inclinations of $82.2\pm1.0^{\circ}$. We estimate the irradiating luminosity to be about $8\times10^{35}$\,erg\,s$^{-1}$, in good agreement with X-ray constraints. | 2S~0921--630 was discovered as an X-ray source by \citet{Li:1978a} using SAS-3 and was identified with an approximate 17th magnitude star, V395~Car. The system shows partial eclipses in both the optical and X-ray bands \citep{Branduardi-Raymont:1983a, Chevalier:1982a, Mason:1987a}. The orbital period is about 9\,days and the system must have an inclination between $70^{\circ}$ and $90^{\circ}$. Optical dips of up to 2 magnitudes deep \citep{Krzeminski:1991a} have been present in optical light curves of V395~Car, but no complete orbital lightcurve has yet been published. The companion star has been identified as a K0\,{\sc iii} star \citep{Shahbaz:1999a}. Several attempts have been made to constrain the system parameters of V395~Car using the photospheric absorption lines from the companion star, which are detectable in spite of the large contamination by disk flux \citep{Shahbaz:1999a,Shahbaz:2004a,Jonker:2005a,Shahbaz:2007a,Steeghs:2007a}. From measurements of the radial velocity curve and rotational broadening, and assuming an appropriate inclination for an eclipsing system, it is possible to solve for the masses of both objects. The remaining uncertainty is in the `K-correction', which accounts for suppression of absorption lines on the inner face of the companion star by irradiation. Assuming negligible K-correction and an inclination of $i=83^{\circ}$, \citet{Steeghs:2007a} deduce a compact object mass of $M_1=1.44\pm0.1$\,M$_{\odot}$, a companion star mass of $M_2=0.35\pm0.03$\,M$_{\odot}$, and hence a mass ratio of $q=0.24\pm0.02$. \citet{Shahbaz:2007a} assume $i=75^{\circ}$ and deduce $M_1 = 1.37\pm0.13$\,M$_{\odot}$ and $q=0.281\pm0.034$. We note that the assumed inclination does not affect the derived mass ratio, and that the two estimates of $q$ are consistent. The mass estimates, too, agree, and both are consistent with a canonical mass neutron star, in contrast to earlier mass estimates \citep{Shahbaz:2004a,Jonker:2005a}, which used the overestimated rotational broadening measurement of \citet{Shahbaz:1999a}. Because V395~Car eclipses, the inclination is relatively well constrained and the derived parameters are relatively insensitive to the remaining inclination uncertainty. Nonetheless, it is still of interest to examine the orbital lightcurve of the binary, and in fact it is an advantage that the system parameters are already well constrained, reducing the uncertainty in the interpretation of the lightcurves. In this paper we present optical and IR photometry from two observing seasons totalling 187 nights of monitoring of V395~Car obtained to investigate both the orbital and longer-term lightcurve. We begin by describing our data in Section~\ref{DataSection}. Section~\ref{LongtermSection} presents the long-term lightcurve, and Section~\ref{PeriodSection} proceeds to derive an improved measurement of the orbital period and ephemeris zero-point based on these data together with archival measurements. With the period firmly established, we present orbital lightcurves in the three bands in Section~\ref{OrbitalSection} and model them in Section~\ref{ModelingSection}. We discuss the implications of our results in Section~\ref{DiscussionSection} and summarise our conclusions in Section~\ref{ConclusionSection}. | \label{ConclusionSection} We have obtained the first full, and multi-colour, orbital lightcurve of the long-period LMXB V395~Car. We refine the orbital period to $9.0026\pm0.0001$\,days and present an updated time of minimum, $T_0 = 2453397.28\pm0.02$ (HJD). We successfully fit our $BVJ$ lightcurves using the XRbinary code and derive a binary inclination of $i=82.2\pm1.0^{\circ}$. This is consistent with previous constraints and does not change previously published compact object mass estimates. We also estimate an X-ray luminosity of $10^{36}$\,erg\,s$^{-1}$, consistent with X-ray estimates of \citet{Kallman:2003a}. This makes V395~Car a relatively low mass-transfer rate system in sharp contrast to 4U~1822--371 where an accretion rate close to the Eddington limit is inferred. \citep{Bayless:2010a}. With the period so close to an integer days, one of the major limitations of our dataset remains the incomplete sampling in orbital phase. Further progress will require a completely sampled lightcurve, and consequently a multi-site campaign spanning a range of longitudes. | 12 | 6 | 1206.0032 |
1206 | 1206.6108_arXiv.txt | We present structural measurements for the galaxies in the $0.05<z<0.0585$ groups of the \emph{Zurich Environmental Study}, aimed at establishing how galaxy properties depend on four environmental parameters: group halo mass $M_{GROUP}$, group-centric distance $R/R_{200}$, ranking into central or satellite, and large-scale structure density $\delta_{LSS}$. Global galaxy structure is quantified both parametrically and non-parametrically. We correct all these measurements for observational biases due to PSF blurring and surface brightness effects as a function of galaxy size, magnitude, steepness of light profile and ellipticity. Structural parameters are derived also for bulges, disks and bars. We use the galaxy bulge-to-total ratios (B/T), together with the calibrated non-parametric structural estimators, to implement a quantitative morphological classification that maximizes purity in the resulting morphological samples. We investigate how the concentration $C$ of satellite galaxies depends on galaxy mass for each Hubble type, and on $M_{GROUP}$, $R/R_{200}$ and $\delta_{LSS}$. At galaxy masses $M\ge10^{10}\Msol$, the concentration of disk satellites increases with increasing stellar mass, separately within each morphological bin of B/T. The known increase in concentration with stellar mass for disk satellites is thus due, at least in part, to an increase in galaxy central stellar density at constant B/T. The correlation between concentration and galaxy stellar mass becomes progressively steeper for later morphological types. The concentration of disk satellites shows a barely significant dependence on $\delta_{LSS}$ or $R/R_{200}$. The strongest environmental effect is found with group mass for $>10^{10}\Msol$ disk-dominated satellites, which are $\sim10\%$ more concentrated in high mass groups than in lower mass groups. | \setcounter{footnote}{0} We present the methodology used to derive structural measurements for the galaxies investigated in the \emph{Zurich Environmental Study (ZENS)} (\citealt{Carollo_et_al_2013}, hereafter Paper I). The resulting measurements are provided in the ZENS global catalog that we have published electronically with Paper I\footnote{The ZENS catalog is also downloadable from: http://www.astro.ethz.ch/research/Projects/ZENS.}. ZENS is designed to address the question of which specific environment is most relevant for influencing the properties of different galaxy populations. Several definitions of environment have been commonly employed in the literature to study the relation between environment and galaxy evolution: the density of galaxies calculated out to a fixed or an adaptive distance \citep[e.g.][]{Dressler_1980,Hogg_et_al_2003, Cooper_et_al_2005,Baldry_et_al_2006}, the mass of the host group or cluster \citep{Weinmann_et_al_2006, Kimm_et_al_2009}, the distance from the group/cluster center \citep{Whitmore_Gilmore_1991,Balogh_et_al_1997,Lewis_et_al_2002, DePropris_et_al_2003,Hansen_2009} or the location into larger structures such as cosmic filaments or superclusters \citep{Einasto_et_al_2007,Porter_et_al_2008}. Recently, the ability to separate galaxies into centrals and satellites within their host group halos has produced mounting evidence that the environmental influence on the star-formation properties of galaxies may peak for satellite galaxies \citep{VanDenBosch_et_al_2008,Peng_et_al_2010,Peng_et_al_2012,Knobel_et_al_2012}. There is cross-talk however between different definitions of environment, which may also relate to one another from a physical perspective. A key question is therefore to identify what is the relative importance of different environmental conditions for well-defined galaxy populations of different masses, star formation activity levels, and structural/morphological properties \citep[e.g.][]{Blanton_Berlind_2007,Wilman_et_al_2010,Peng_et_al_2012,Muldrew_et_al_2012,Woo_et_al_2012}. In ZENS we aim at helping clarifying which of the many environments that a galaxy experience has a larger impact on its evolution. We do so by using the same sample of suitably selected nearby galaxies to investigate the dependence of their properties, at fixed stellar mass, on four environment measurements: the host group halo mass, the radial segregation within the group, the large scale density field on which the group halos reside, and the galaxy rank within its group halo, i.e., whether it is the central or a satellite galaxy within the gravitational potential of its host group. In computing our proxies for these different environments, we have attempted to minimize cross-talks between their definitions, in order to better disentangle one from another of the physical conditions that galaxies experience (see Paper I). This paper focuses on the quantification of robust galaxy structural and morphological properties, which provide key information on the life histories of galaxies. The presence and properties of massive disks and spheroids highlight the occurrence of relatively slow and dissipative gas accretion \citep{White_Rees_1978,Fall_Efstathiou_1980} or mergers \citep[e.g][]{Toomre_1977,Barnes_1988,Schweizer_1990,Naab_Burkert_2003}, respectively. Inner cores or cusps \citep{Ferrarese_et_al_1994,Lauer_et_al_1995,Carollo_et_al_1997a,Carollo_et_al_1997b,Graham_Guzman_2003,Graham_et_al_2003,Truijillo_et_al_2004,Cote_et_al_2007, Kormendy_et_al_2009} and tidal debris \citep{Malin_Carter_1980,Malin_Carter_1983,Forbes_et_al_1992,vanDokkum_2005, Tal_et_al_2009,Janowiecki_et_al_2010} also trace the degree of dissipation involved in the evolution of galaxies (not surprisingly with some debate, \citealt[][]{Mihos_Hernquist_1994, Mihos_Hernquist_1996, Kawata_et_al_2006,Feldmann_et_al_2008,Hopkins_et_al_2009}). Bars and pseudo-bulges are smoking guns for either secular evolution processes (\citealt{Kormendy_1979, Combes_et_al_1990,Courteau_et_al_1996,Norman_et_al_1996,Wyse_et_al_1997,Carollo_1999, Carollo_et_al_1997c,Carollo_et_al_1998,Carollo_et_al_2001,Balcells_et_al_2003,Debattista_et_al_2004,Debattista_et_al_2006,Kormendy_Kennicutt_2004,Fisher_Drory_2008}), or possibly for early bulge formation through instabilities in the proto-disks \citep{Immeli_et_al_2004,Carollo_et_al_2007,Dekel_et_al_2009a,Dekel_et_al_2009b}. A robust determination of galaxy morphology is hence essential for understanding whether this is linked to any of the environmental conditions above, how precisely it relates to the occurrence, enhancement or cessation of star formation activity, and thus for pinning down which physical processes drive galaxy evolution. Determining galaxy structure is however notoriously not a trivial task. Visual morphological classification is still a widely adopted method \citep[e.g.][]{Lintott_et_al_2008,Nair_Abraham_2010}, despite its subjectivity and failure to provide quantitative measurements for different components, which are necessary to trace galaxy assembly over cosmic time. For these reasons numerous publications have been devoted to the development of methods and software for the automated quantification of structure on large galaxy samples. There are a number of approaches to the problem which can be broadly divided into two categories: those which employ parametric descriptors for the galaxy morphology, namely a set of analytical profiles used to model the bulge, disk, or bar component \citep[e.g.][]{Simard_et_al_2002, Peng_et_al_2002,deSouza_et_al_2004} and those which instead use the observed properties of the light distribution, such as the degree of asymmetry, isolation of bright pixels or decompositions into a set of basis functions \citep[e.g.][]{Abraham_et_al_1996,Conselice_2003,Refregier_2003,Lotz_et_al_2004,Scarlata_et_al_2007}. The two methods have different strengths: parametric decomposition is useful to obtain measurements of characteristic sizes and to have an estimate of the relative importance of the bar, disk and bulge components; it also easily includes the effects of seeing. Non-parametric estimators well describe the inhomogeneities in the light distributions of real galaxies, which typically display irregular, non axis-symmetric features generated by recent star-formation, dust or galaxy interactions. Both parametric and non-parametric measurements suffer however from a number of observational biases, which must be corrected for in order to perform comparisons between galaxies of different properties, and observed in different conditions. In particular, in ground-based surveys, the effect of atmospheric seeing is one of the major complications. Several studies have shown the strong impact of the seeing on the photometric and structural properties of galaxies, not only in the inner regions of galaxies, but also out to radii corresponding to several FWHM of the Point Spread Function (PSF) \citep[e.g.][]{Schweizer_1979,Franx_et_al_1989, Saglia_et_al_1993,Trujillo_et_al_2001,Graham_2001}. Another factor which affects the derivation of structural parameters is the inclination angle at which a galaxy is observed. The overlap, in projection, of multiple subcomponents, as well as physical factors such as the non-uniform distribution of inter-stellar dust -- which causes a higher attenuation of short-wavelength light in the central regions of edge-on galaxies than in similar face-on galaxies \citep[e.g.][]{Driver_et_al_2007,Shao_et_al_2007} -- can substantially bias the measurements of sizes, bulge-to-disk ratios, concentration and even stellar masses \citep[e.g.,][]{Maller_et_al_2009,Graham_Worley_2008, Bailin_Harris_2008}. Finally, the background sky makes the detection of faint components difficult, a fact which introduces severe biases in the measurements of the galaxy properties, especially magnitudes and sizes \citep[e.g.][]{Disney_1976,Impey_Bothun_1997}. The strength of the bias depends on galaxy size, inclination and stellar light profile. Although a number of widely-used measurement techniques, such as the computation of Kron aperture fluxes \citep{Kron_1980} or the extrapolation of model galaxy surface brightness profiles \citep{Sersic_1968}, can help recover light below the isophotal limit, significant systematic biases remain in the low surface brightness regimes specific to each survey \citep[e.g.][]{Graham_et_al_2005, Cameron_Driver_2007,Cameron_Driver_2009, Haussler_et_al_2007}. In ZENS we attempt {\it to correct}, when possible, all measurements of galaxy structure for systematic biases as a function of PSF-size, and also galaxy magnitude, size, axis ratio and radial shape of the light profile. We also quantify the size of systematic biases in regimes of parameter space where the (statistical) recovery of the intrinsic information is not achievable, e.g., at small galaxy sizes and low surface brightnesses. We also stress that in ZENS each galaxy is handled individually, till self-consistent and both quantitatively- and visually-checked accurate measurements are achieved. This enables ZENS to tackle complementary questions regarding the galaxy-environments relationship relative to larger but less detailed galaxy samples. The paper is organized as follows. In Section \ref{sec:survey} we briefly review the specifications and definitions for the four environments under scrutiny in ZENS. We then devote the first part of the paper to an overview of the structural measurements carried out on the ZENS galaxy sample. These measurements include isophotal analyses and bar detection/quantification (Section \ref{sec:IsophotalAnalysis}), analytical surface brightness fits and bulge+disk decompositions (Section \ref{sec:GIM2Dmeasurements}), and derivation of non-parametric structural indices (concentration, Gini, asymmetry, M$_{20}$, smoothness; Section \ref{sec:ZEST}). In Section \ref{sec:Simulations} we thoroughly investigate the sources of error in these measurements and derive a correction scheme that recovers the intrinsic structural parameters. In Section \ref{sec:MorphClass} we present the morphological classification of the \textsc{ZENS} galaxies, based on a quantitative partition of the structural parameter space in regions that are associated with elliptical, bulge-dominated disks, intermediate bulge-to-total ratio disks, late-type disks and irregular galaxies. In the same Section we also discuss the statistics of the structural properties for the various morphological classes. We describe in detail in Appendix \ref{app:DataReduction} the data reduction and photometric calibration of the ESO $B$ and $I$ WFI/2.2m imaging data for the ZENS groups that are introduced in Paper I. Appendices \ref{sec:testGIM2D}-\ref{app:corrections_PSF} present additional details on the tests performed on the analytical surface-brightness fits and supplementary information for the derivation of the corrections for the structural parameters. Stamp images for galaxies in the different morphological classes are found in Appendix \ref{app:ClassDist}. In the final part of the paper we use the corrected structural measurements to study, at constant stellar mass, the concentration of satellite galaxies as a function of Hubble type and environment (Section \ref{sec:Results}). We summarize the paper in Section \ref{sec:Conclusions}. This second \textsc{ZENS} publication is complemented by a companion paper (\citealt{Cibinel_et_al_2013}, hereafter Paper III), in which we present the spectrophotometric properties of our galaxy sample and to which we refer for details on the derivation of, e.g., the galaxy spectrophotometric types, star formation rates and stellar masses that we use in this and other ZENS papers. The following cosmological parameters are adopted in all the ZENS publications: $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ and $h=0.7$. Unless otherwise specified, magnitudes are in the AB system, and galaxy sizes are \emph{semi-major} axis measurements. All derived luminosities are corrected for Galactic extinction using the maps of \citealt{Schlegel_et_al_1998}. | \label{sec:Conclusions} We have presented detailed structural analyses performed on the 1455 galaxies in the 141 ZENS groups introduced in Paper I. We remark that the corresponding ZENS catalogue has however 1484 galaxy with valid entries, since we measure parameters for individual galaxies that are members of 29 galaxy pairs, which are given as one single entry in the parent 2dFGRS galaxy catalogue. The parametric and non-parametric structural measurements presented here, together with the detailed environmental parameters from Paper I and the photometric (including stellar masses) measurements presented in Paper III, set the basis for a number of forthcoming publications which use the \textsc{ZENS} data to explore which environmental scales are relevant for the evolution of morphologically different galaxy populations at different mass scales. The measurements are published in the global ZENS catalog with Paper I. In detail, the measurements that we have presented here are: \begin{itemize} \item Strength of bars in disks, quantified through an isophotal analysis (Section \ref{sec:IsophotalAnalysis}); \item Single- and double-component (bulge+disk) S\'ersic fits parameters both in the $B$ and $I$ bands, including model-based galaxy sizes (Sections \ref{sec:SersicFits} and \ref{sec:GIM2D_decomp}), as well as bulge+disk+bar fits for disks with a noticeable bar component (Appendix \ref{app:BulgeBarDisk}); \item Non-parametric structural indices of concentration, asymmetry, smoothness of the light distribution, and Gini and $M_{20}$ coefficients, and ``aperture photometric" size estimates (Section \ref{sec:ZEST}); \item Morphological classes, based on a quantitative bulge+disk criterion, augmented (or supported) by quantitative criteria regarding the non-parametric diagnostics, after correction for observational biases (Sections \ref{sec:GIM2D_decomp} and \ref{sec:MorphClass}). \end{itemize} Crucially, we do indeed derive correction matrices, which we apply to the relevant structural estimates, to minimize biases which, depending on PSF size as well as galaxy magnitude, size, shape of light profile and ellipticity, would otherwise prevent a reliable comparison of the structural properties of galaxies observed in different seeing conditions, and lying in different regions of this four-dimensional galaxy parameter space (see Section \ref{sec:Simulations}). As expected, biases in the model-fit parameters are substantially reduced thanks to the treatment of the PSF-blurring effects in these algorithms; still, some are present even in the model-fit parameters, which may have an impact in some analyses if left uncorrected. Disk properties are well measured and require very modest or no further corrections; so are bulge-to-total ratios, which therefore offer an excellent parameter to base a quantitative morphological/structural classification. In contrast, bulge $n$ S\'ersic indices and half-light radii are degenerate in some circumstances. This implies that, on global galactic scales, concentration parameter or n-S\'ersic index, alone, are not a good proxy for bulge-to-disk ratio and thus morphology: bulges can give large contributions to the total light budget and have large half-light radii, leading to low galaxy concentrations, or, vice-versa, bulges can be very compact and lead to high galaxy concentrations despite a modest contribution to the total light. Finally, we particularly warn against using uncorrected non-parametric estimators as galaxy classifiers since, understandably, they suffer from severe observational biases and introduce severe errors in the classifications (see Section \ref{sec:Simulations}). As a first application of our corrected structural measurements, we have studied the variation of concentration in \emph{satellite} galaxies of fixed stellar mass (from Paper III) with morphological type, and with the three environments detailed in Paper I, i.e., the mass of the host group halo, the projected group-centric distance, and the density of the large-scale structure cosmic web (Section \ref{sec:Results}). We find that the known correlation of satellite concentration with galaxy stellar mass \citep[e.g.][]{Kauffmann_et_al_2003} holds at a fixed morphological type. Specifically, there is a genuine increase in concentration with increasing stellar mass for disk satellite galaxies within each separate bin of bulge-dominated galaxies, intermediate-type and late-type disk morphologies. The slope of the concentration vs.\, galaxy stellar mass relationship flattens from the later to the earlier types (and becomes $\sim0$ for satellites with an elliptical morphology, which however cover a limited range in galaxy stellar mass in our sample). It is not trivial to disentangle, in a physically meaningful way, the contributions of an increasing bulge-to-total ratio with increasing mass, and an increasing bulge concentration at fixed bulge-to-total ratio with increasing mass, to the increase in concentration of satellite disk galaxies with increasing stellar mass. This is true even in our study, in which we did indeed base our morphological classification on the bulge-to-total ratio, and thus bulge-to-total variations within each Hubble type should be minimized. There are nevertheless some residual effects, as the median $B/T$ is found to vary in our sample from $\sim7\%$ at $10^{10}\Msol$ to $\sim12\%$ at $\sim10^{10.5}\Msol$ for late-type disks, and from 59$\%$ to 64$\%$ between $10^{10}\Msol$ and $10^{11}\Msol$ for bulge-dominated disks. Still, for intermediate-type disk satellites, in which the bulge-to-total ratio is tightly constrained by firm lower an upper boundaries by definition, the increase in satellite density with stellar mass can be genuinely ascribed to a $\sim 30\%$ decrease in the bulge-to-disk size ratio, i.e., to a genuine increase in the concentration of the bulge component at fixed bulge-to-total light ratio. We tentatively assume this as the explanation for the increase of disk satellite concentration with stellar mass at fixed Hubble type, an hypothesis which we will test with further analyses. When considering the galaxies environment, we find that, at galaxy stellar mass $\sim10^{10} \Msol$ and above, $(i)$ bulge-dominated satellites tend to be marginally more concentrated at low LSS densities and high group-centric distances, and $(ii)$ in contrast, disk-dominated satellites are significantly more concentrated in high group masses. The interpretation of the first weak trend as a hint for an environmental effect is further hampered, at $M\sim10^{11}\Msol$, by a higher fraction of S0s relative to bulge-dominated spiral galaxies in low $\delta_{LSS}$ relative to high densities (60\% and 0\%, respectively in $M_{GROUP}<10^{13.5} \Msol$ groups and 58\% and 35\% over all groups) and in the outskirts relative to the cores of groups (50\% and 31\%, over all groups). This change in morphological mix at the high-end of the B/T sequence is consistent with other observational works which have also reported an increase in the S0 fraction in the outskirts of groups and clusters, although with no distinction between central and satellite galaxies. The early study of \cite{Whitmore_et_al_1993} on local cluster reveals a drop in the S0 fraction close to the cluster centers, which they interpret as the outcome of disk galaxies destruction happening at the cluster cores. The more recent analysis of the morphology-density relation at redshift $0.05<z<0.1$ by \citealt{Goto_et_al_2003} shows a depletion of S0 galaxies within 0.3 virial radii, and an increase in the early spirals population at the same distances, consistent with the variation we see in our ZENS sample. Also in $z\sim0.4$ groups \citealt{Wilman_et_al_2009} find hints for an excess of S0 galaxies at $r\gtrsim0.3$ Mpc with respect to the group centers, which instead host a higher fraction of late-type disks. These authors furthermore find that the fraction of S0 galaxies in groups is comparable to the one in clusters at the same redshift, suggesting that galaxy pre-processing and S0 formation is effective already at these low densities. The increase of S0 in the outer regions of the groups is particularly interesting, as it counter to the intuitive idea that gas-rich bulge-dominated spirals may become S0 galaxies as they fall deeper into their group potential wells (see e.g. \citealt{Bekki_et_al_2002} who find that gas stripping in the group environment is only effective close to the group centers). Speculating on the above, a possible scenario is that $(i)$ bulge-dominated galaxies are transformed into S0s as soon as disk galaxies enter the group potential, and $(ii)$ further stellar evolution within the groups replenished the dried-out disks of S0 galaxies of fresh gas, establishing/restoring in them a bulge-dominated morphology. Note that an increase in the fraction of dusty star-forming galaxies at high densities and close to cluster centers has been observed at low redshift, e.g. \citealt{Gallazzi_et_al_2009,Mahajan_Raychaudhury_2009}. Furthermore, the evidence of polar and extended HI disks around S0s -- e.g. \citealt{vanGorkom_1987,Noordermeer_et_al_2005,Sage_Welch_2006} -- may also support this "disk-regrowth" scenario. The second, statistically more significant trend for $M>10^{10} M_\odot$ disk-dominated galaxies to have a higher concentration in high-mass ($M>10^{13.5} M_\odot$) than in lower-mass groups appears to be a genuine environmental effect which should suffer from no morphological complications. \citet{Weinmann_et_al_2009,Guo_et_al_2009} found evidence in their SDSS sample that, among galaxies with $C<3$, satellites are more concentrated than centrals with identical stellar mass. These authors interpret the variation in concentration within the framework of gradual stripping and subsequent quenching of satellite galaxies during infall into the group potential. This is also believed to cause a reddening of the satellite galaxies and a shrinking of their typical sizes. While we postpone an analysis on the subject to a forthcoming paper, here we wish to note that the SDSS studies divide the early- and late-type morphological classes on the basis of a non-corrected concentration criterion ($C<3$ for late-type and $C>3$ for early-type galaxies). Comparing their Figure 1 with our Figure \ref{fig:C_vs_BT}, we see that $C<3$ in the SDSS system roughly corresponds to a value of $C=3.5-4$ in ZENS, as our concentrations show a small offset due to the corrections we applied. As illustrated in Figure \ref{fig:Cmass}, a cut at constant concentration sub-divides the \textsc{ZENS} galaxy sample into two broad bins in which, however, individual Hubble types are mixed together below the chosen $C$ threshold. It thus remains an open question whether the difference in concentration between centrals and satellites reported by those authors is the result of a variation in the morphological mix of the central vs.\, satellite populations, or rather a change in the structure of central and satellite galaxies at a fixed morphological type. We will address this question in a dedicated ZENS analysis. | 12 | 6 | 1206.6108 |
1206 | 1206.6614_arXiv.txt | \snr\ is a young Galactic supernova remnant (SNR), whose identification as the remains of a Type-II supernova (SN) explosion has been debated for a long time. In particular, recent multi-wavelength observations suggest that it is the result of a Type Ia SN, based on spectroscopy of the SNR shell and the lack of a compact stellar remnant. However, two X-ray sources, one detected by \ein\ and \ros\ (Source V) and the other by \chan\ (Source N) have been proposed as possible isolated neutron star candidates. In both cases, no clear optical identification was available and, therefore, we performed an optical and X-ray study to determine the nature of these two sources. Based on \chan\ astrometry, Source V is associated with a bright $V\sim14$ star, which had been suggested based on the less accurate \ros\ position. Similarly, from \vltn\ (\vlt) archival observations, we found that Source N is associated with a relatively bright star ($V=20.14 $). These likely identifications suggest that both X-ray sources cannot be isolated neutron stars. | The Galactic supernova remnant (SNR) \snr\ is a very young object, supposed to be the remains of the historical supernova SN 185 A.D. (Clark \& Stephenson 1975). The remnant distance is uncertain, with values ranging from 1 to 2.8 kpc (Kaastra et al.\ 1992; Rosado et al.\ 1996; Sollerman et al.\ 2003). \snr\ has been observed in radio (Hill 1967), in the soft X-rays by \ros\ (e.g., Vink et al.\ 2000), and hard X-rays by {\em ASCA} (e.g, Borkowski et al.\ 2001) and, more recently, both by \xmm\ and \chan\ (Vink et al.\ 2006). In the optical, the southwest edge of \snr\ is associated with the RCW\, 86 nebula (e.g., Smith 1997). At higher energies, it has been detected at TeV $\gamma$-rays by {\em H.E.S.S.} (Aharonian et al.\ 2009). Its morphology is typical of a young shell-type SNR, but with a marked asymmetry in the X-ray surface brightness profile between the fainter northeast and the brighter southwest edges (Vink et al.\ 2000), as observed both in radio (Dickel et al.\ 2001) and in the optical, where only the region associated with the RCW\, 86 nebula is visible (Smith et al.\ 1997). This might imply an off-centred cavity SN explosion, with its centre closer to the southwest edge of the SNR and near the centre of the RCW\, 86 nebula. Recently, Williams et al.\ (2011) suggested that \snr\ is a Type-Ia SNR, based on optical and X-ray spectroscopy of the remnant and the lack of a compact source at the centre, which could be identified as the isolated neutron star (INS) formed out of a Type II SN explosion. In the radio band, Kaspi et al.\ (1996) could not identify a radio pulsar within the SN shell. In the X-rays, Vink et al.\ (2000) spotted an unresolved source (hereafter Source V) $\sim 7\arcmin$ southwest of the SNR centre, in the direction of the RCW\, 86 nebula. However, the possible association with a $V\sim 14$ star coincident with the \ros\ 5\arcsec\ error circle and evidence for long-term variability suggested that this X-ray source was not an INS, although only a better localisation from X-ray observation at higher spatial resolution can either confirm or refute the association. More recently, Gvaramadze \& Vikhlinin (2003) found two X-ray sources with \chan, located on the X-ray bright southwest edge of the \snr\ SNR. Only the northernmost one (dubbed Source N) had no optical counterpart ($R_F\ga 21$) and was proposed as the stellar remnant of the SN explosion. However, its possible identification as an INS has never been confirmed so far. Source N has not been detected in radio. One possibility would be that it is a radio-silent INS, possibly a Compact Central Object, or CCO (De Luca\ 2008). A search for a stellar remnant in the central part of the SNR was carried out with \chan\ by Kaplan et al.\ (2004), who identified all the detected X-ray sources either as foreground stars or background AGNs. Thus, Source V and N remain the only unidentified X-ray sources potentially associated with the remnant of the SN explosion. We investigated their nature using archival \chan\ data to better determine their positions, verify the association of Source V with the bright nearby star, and search for Source N's optical counterpart in archival \vltn\ (\vlt) data. | To establish the nature of Source N, we evaluated whether the flux of its optical counterpart is compatible with what expected for an INS. For the estimated age of \snr\ ($\sim 1800$ yrs) and an optical luminosity comparable to the Crab pulsar, we would expect a flux of $V\sim 16.6$--21.1, after re-normalising for the \snr\ distance (1--2.8 kpc) and interstellar extinction $A_V\sim2.8$--5, inferred upon the N$_{\rm H}$ derived from the PL fit to the \chan\ spectrum using the relation of Predhel \& Schmitt (1995). These values are compatible with the flux of Source N's counterpart ($V=20.4$), at least for the highest values of the interstellar extinction and distance. Alternatively, Source N might be a CCO. For a rotational energy loss rate $\dot{E}<10^{36}$ erg s$^{-1}$ ( see De Luca et al.\ 2012 and references therein), the expected optical luminosity would then be at least two orders of magnitude lower than the Crab pulsar, yielding a flux $V\ga23.5$. We also verified whether the X-ray--to--optical flux ratio of Source N is compatible with that of an INSs. According to the PL model, the observed flux in the 0.3--8 keV energy range is $4.2_{-0.5}^{+0.2}\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$, corresponding to an unabsorbed flux $F_{\rm X} = 8.5_{-0.9}^{+0.4}\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$. The unabsorbed optical flux of Source N's optical counterpart, computed upon the corresponding N$_{\rm H}$, is $F_{\rm opt} = (0.28-2.1) \times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$. Thus, the unabsorbed X-ray--to--optical flux ratio of Source N would be $F_{\rm X}/F_{\rm opt} \sim 0.036$--0.32, by far lower than expected for an INS, for which this ratio is usually of 1000, or higher (e.g., Mignani 2011). Assuming other best-fitting X-ray spectral models also gives very modest values of the $F_{\rm X}/F_{\rm opt}$ ratio. This implies that Source N must be a different type of X-ray source, maybe an AGN, which would be compatible with its $F_{\rm X}/F_{\rm opt}$ ratio (e.g. la Palombara et al.\ 2006) and its possible PL spectrum. Spectroscopy of its optical counterpart is needed to determine the nature of Source N. Thus, from their associations with bright optical counterparts, we conclude that neither Source V nor Source N are INSs. No other unidentified point-like X-ray sources which can be considered possible INS candidates have been discovered by \chan\ surveys of \snr. The available \chan\ observations (Fig.\ 1, top left) cover a large fraction of the SNR. In particular, the most likely locations where the compact remnant of the SN explosion is expected to be found, i.e. the central regions of the SNR and the RCW\, 86 nebula, have been deeply scrutinised (e.g., Kaplan et al.\ 2004; Gvaramadze \& Vikhlinin 2003). Thus, it is unlikely that potential INS candidates were missed. We conclude that the lack of an identified INS supports the conclusion that \snr\ was not born after a Type II SN explosion (Williams et al.\ 2011). | 12 | 6 | 1206.6614 |
1206 | 1206.3494_arXiv.txt | The new results of our observing campaign targeting the isolated neutron star 2XMM~J104608.7-594306 in the Carina Nebula are used to understand how peculiar groups of isolated neutron stars relate to each other, as well as to the bulk of the normal radio pulsar population. | 12 | 6 | 1206.3494 |
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1206 | 1206.3946_arXiv.txt | Gamma-ray bursts (GRBs) that emit photons at GeV energies form a small but significant population of GRBs. However, the number of GRBs whose GeV-emitting period is simultaneously observed in X-rays remains small. We report \gr~observations of \grb~using Fermi's Large Area Telescope (LAT) in the energy range 100~MeV--20~GeV. Gamma-ray emission at these energies was clearly detected using data taken between 180~s and 580~s after the burst, an epoch after the prompt emission phase. The GeV light curve differs from a simple power-law decay, and probably consists of two emission periods. Simultaneous Swift/XRT observations did not show flaring behaviors as in the case of GRB~100728A. We discuss the possibility that the GeV emission is the synchrotron self-Compton radiation of underlying ultraviolet flares. | Since the launch of the Fermi satellite in 2008, more than 20 \gr~bursts (GRBs) have been detected above $\sim$100~MeV by the Large Area Telescope (LAT) aboard the satellite~\citep{lat_090510,lat_080825c,lat_080916c,lat_090902b}. Long-lived MeV--GeV emission of GRBs, first detected in the EGRET era, is now a common feature of LAT-detected GRBs. The nature of such temporally extended emission beyond the prompt GRB phase is not well understood. A widely-discussed radiation mechanism is synchrotron emission from external shocked electrons~\citep[e.g.,][]{Zou09,BarniolDuran09,Ghisellini10}, but inverse-Compton scattering off flare photons or late-time activities of the central engine are among alternative scenarios~\citep{lat_100728a,Zhang11}. Simultaneous observations at other wavelengths of such extended MeV--GeV emission from GRBs are crucial to disentangle various emission models. By May 2011, only two LAT-detected GRBs have been simultaneously observed by \emph{Swift}'s X-ray telescope (XRT) during its GeV-emitting period: GRB~090510 and GRB~100728A. GRB~090510 remains the only short GRB detected by LAT, and its GeV emission can be interpreted as, e.g., synchrotron radiation of the forward shock electrons~\citep[][but see \citet{Gao2009} and \citet{Piran_Nakar10} for an opposing viewpoint]{090510_afterglow,Ghirlanda_090510}. In the case of GRB~100728A, X-ray flares and corresponding GeV emission were detected up to $\sim$1~ks after the burst, suggesting their common origin~\citep{lat_100728a}. However, the number of GRBs whose GeV emission is simultaneously covered in X-rays remains low. In this paper, we report another such case: \grb, that was detected by Fermi/LAT and Swift/XRT simultaneously for several hundred seconds. Errors are reported at 1$\sigma$, unless otherwise specified. | \grb~is the third GRB detected by Fermi/LAT and Swift/XRT simultaneously. We have shown that the GeV light curve differs from a simple power-law decay, and probably consists of two emission periods. The rapid decrease of GeV flux during both periods challenges the notion that the emission comes from the external forward shock. While in the case of GRB~100728A, late-time X-ray flares seem to accompany the GeV emission, no such flares are seen in the time frame during which GeV emission was detected. This suggests a different origin of the GeV emission between the two cases. We discuss the possibility that the GeV emission is the SSC radiation of an underlying ultraviolet (UV) flare. Multiwavelength coverage of the rare class of LAT GRBs during the GeV-emitting period is crucial. | 12 | 6 | 1206.3946 |
1206 | 1206.6564_arXiv.txt | Kinetic plasma theory is used to generate synthetic spacecraft data to analyze and interpret the compressible fluctuations in the inertial range of solar wind turbulence. The kinetic counterparts of the three familiar linear MHD wave modes---the fast, Alfv\'en, and slow waves---are identified and the properties of the density-parallel magnetic field correlation for these kinetic wave modes is presented. The construction of synthetic spacecraft data, based on the quasi-linear premise---that some characteristics of magnetized plasma turbulence can be usefully modeled as a collection of randomly phased, linear wave modes---is described in detail. Theoretical predictions of the density-parallel magnetic field correlation based on MHD and Vlasov-Maxwell linear eigenfunctions are presented and compared to the observational determination of this correlation based on 10 years of\emph{Wind} spacecraft data. It is demonstrated that MHD theory is inadequate to describe the compressible turbulent fluctuations and that the observed density-parallel magnetic field correlation is consistent with a statistically negligible kinetic fast wave energy contribution for the large sample used in this study. A model of the solar wind inertial range fluctuations is proposed comprised of a mixture of a critically balanced distribution of incompressible Alfv\'enic fluctuations and a critically balanced or more anisotropic than critical balance distribution of compressible slow wave fluctuations. These results imply that there is little or no transfer of large scale turbulent energy through the inertial range down to whistler waves at small scales. | Despite more than forty years of direct spacecraft measurements of turbulence in the near-Earth solar wind \citep{Coleman:1968}, our understanding of turbulence in a magnetized plasma remains incomplete. One of the primary goals is to understand the role of the turbulence in mediating the transfer of energy from large to small scales. Within the turbulent inertial range of scales, corresponding to spacecraft-frame frequencies of $10^{-4}$ Hz $\lesssim f_{sc} \lesssim 1$ Hz or length scales $10^6$~km $\gtrsim \lambda \gtrsim 10^2$~km, the fluctuations involved in this energy transfer are a mixture of compressible and incompressible fluctuations, with around 90\% of the energy in the incompressible component \citep{Tu:1995,Bruno:2005}. These incompressible fluctuations have been identifed as \Alfven waves \citep{Belcher:1971}, but the nature of the compressible component remains uncertain. The compressible turbulent fluctuations have often been interpreted as a combination of magnetoacoustic (fast MHD) waves and pressure-balanced structures \citep{Tu:1995,Bruno:2005}. Early studies of thermal and magnetic pressure fluctuations in the solar wind found an anti-correlation of the thermal pressure and magnetic pressure at timescales of 1~h, corresponding to an interval of constant total pressure, or a pressure-balanced structure (PBS) \citep{Burlaga:1968,Burlaga:1970}. Further studies found evidence of PBSs out to 24~AU \citep{Burlaga:1990}. Related investigations discovered a similar anti-correlation between the density $n$ and magnetic field magnitude $B$ from 0.3~AU to 18~AU on timescales ranging from several hours to 1.8 minutes \citep{Vellante:1987,Roberts:1987b,Roberts:1987a,Roberts:1990}. Theoretical studies of compressive MHD fluctuations in the low-Mach number, high-$\beta$ limit interpreted these anti-correlated density-magnetic field strength observations as nonpropagating ``pseudosound'' density fluctuations \citep{Montgomery:1987,Matthaeus:1991}. Later, a more comprehensive observational investigation confirmed the general density-magnetic field magnitude anti-correlation, but also identified a few positively correlated intervals consistent with the magnetosonic (fast MHD) wave \citep{Tu:1994}. Analysis of \emph{Ulysses} observations found evidence for PBSs at inertial range scales in the high latitude solar wind \citep{McComas:1995,Reisenfeld:1999,Bavassano:2004}. Studies of the electron density up to $f=2.5$~Hz also found pressure balanced structures but interpreted these as ion acoustic (slow MHD) waves and recognized that PBSs are simply the ion acoustic (slow MHD) wave in the perpendicular wavevector limit \citep{Kellogg:2005}, a fact previously noted by \citet{Tu:1994}. Recently, measurements of the anti-correlation between electron density and magnetic field strength indicated the existence of PBSs over timescales ranging from 1000~s down to 10~s \citep{Yao:2011}. An important consideration in the study of the compressible fluctuations of solar wind turbulence is the fact that the mean free path in the solar wind plasma is about 1~AU, so the dynamics over the entire inertial range is weakly collisional. The implications of this fact have not been seriously addressed in any of the aforementioned studies of compressible fluctuations in solar wind turbulence. The MHD description is rigorously valid only in the limit of strong collisionality, so a kinetic description is formally required to describe the inertial range turbulence. In the limit $k_\parallel \ll k_\perp$ predicted by anisotropic MHD turbulence theories\footnote{Parallel and perpendicular are defined with respect to the direction of the local mean magnetic field.} \citep{Goldreich:1995,Boldyrev:2006}, it has been demonstrated that, even in the weakly collisional limit, the turbulent dynamics of the \Alfven waves decouples from the compressible fluctuations and is rigorously described by the equations of reduced MHD \citep{Schekochihin:2009}. The compressible fast and slow wave modes, on the other hand, require a kinetic description to resolve both the wave dynamics and the collisionless kinetic damping mechanisms. The study presented here is the first to examine the properties of the compressible fluctuations in the turbulent solar wind using Vlasov-Maxwell kinetic theory. Specifically, we use the predicted correlation between the density fluctuations and parallel magnetic field fluctuations to determine the nature of the compressible fluctuations in the solar wind. In \secref{sec:collisionless}, we explore the connection between the familiar linear wave modes in MHD and the corresponding kinetic wave modes in Vlasov-Maxwell kinetic theory and demonstrate that weakly collisional conditions do not change the qualitative properties of the density-parallel magnetic field correlation. In \secref{sec:ssd}, we discuss the quasi-linear premise upon which the method of synthetic spacecraft data is based and describe in detail the procedure for generating synthetic spacecraft data. The synthetic spacecraft data predictions of the density-parallel magnetic field correlation based on linear MHD and Vlasov-Maxwell eigenfunctions is presented in \secref{sec:ssd_predict}. A comparison of the synthetic spacecraft data predictions to the observational determination of the density-parallel magnetic field correlation is presented in \secref{sec:discuss}, showing a statistically negligible fast wave energy contribution to the compressible fluctuations. The implications of this finding are discussed before summarizing the findings of this investigation in \secref{sec:conc}. | \begin{enumerate} \item The predicted $C(\delta n, \delta B_{\parallel})$ and its dependence on plasma $\beta_i$ using linear MHD eigenfunctions is significantly different from the prediction using linear Vlasov-Maxwell eigenfunctions. Only the prediction based on kinetic theory appears to agree with the spacecraft measurements, leading to the expected conclusion that MHD theory is inadequate to describe the compressible fluctuations in the weakly collisional solar wind. \item Strong \emph{a posteriori} evidence for the validity of the quasi-linear premise is provided by the striking agreement between the observationally determined $C(\delta n, \delta B_{\parallel})$ over a very large statistical sample and the predicted $C(\delta n, \delta B_{\parallel})$ based on synthetic spacecraft data. \item The observed $C(\delta n, \delta B_{\parallel})$ computed in a companion work \citep{Howes:2012a} is consistent with a statistically negligible kinetic fast wave energy contribution for the large sample used in this study. Note, however, that the our companion work also found that a very small fraction of the intervals have $C(\delta n ,\delta B_{\parallel})>0$, possibly indicating a trace population of fast waves \citep{Howes:2012a}. \item The quantitative dependence of $C(\delta n, \delta B_{\parallel})$ on the ion plasma beta $\beta_i$ provides evidence that the slow wave fluctuations are not isotropically distributed, but rather have an anisotropic distribution, that is possibly given by the condition of critical balance \citep{Goldreich:1995} or that is more anisotropic than critical balance. \end{enumerate} In conclusion, our analysis using kinetic theory to interpret the compressible fluctuations motivates the following physical model of the turbulent fluctuations in the solar wind inertial range. In this model, the solar wind inertial range consists of a mixture of turbulent fluctuations, with 90\% of the energy due to incompressible \Alfvenic fluctuations and the remaining 10\% of the energy due to compressible slow wave fluctuations. The \Alfvenic turbulent power is distributed anisotropically in wave vector space according to critical balance. The turbulent \Alfven wave dynamics advects and cascades the slow wave fluctuations to smaller scales at the \Alfven wave frequency. The slow wave turbulent power may be either critically balanced, or more anisotropic than critical balance due to collisionless damping of the slow wave fluctuations. Since only the fast wave turbulent cascade is expected to nonlinearly transfer energy to whistler waves at $k \rho_i \gtrsim 1$, and the frequency mismatch between the fast and either the \Alfven or slow waves should prevent non-linear coupling, there is little or no transfer of large scale turbulent energy through the inertial range down to whistler waves at small scales. Therefore, any whistler wave fluctuations at scales $k \rho_i \gtrsim 1$ must be generated by some other process, \emph{e.g.}, kinetic temperature anisotropy instabilities \citep{Kasper:2002,Hellinger:2006,Bale:2009} or kinetic drift instabilities driven by differential flow between protons and alpha particles \citep{McKenzie:1993,Kasper:2008,Bourouaine:2011}. Finally, the lack of statistically significant fast wave energy has important implications for efficient numerical modeling of solar wind turbulent fluctuations. This work demonstrates clearly the importance of a kinetic approach to model adequately the turbulent fluctuations, yet a general kinetic numerical treatment---\emph{e.g.}, the Particle-In-Cell (PIC) method---in three spatial dimensions (required for physically relevant modeling of the dominant nonlinear interactions in solar wind turbulence \citep{Howes:2011a}) is too computationally costly to be presently feasible. Fortunately, it is possible to perform kinetic numerical simulations of solar wind turbulence in three spatial dimensions using gyrokinetics, a rigorous, low-frequency, anisotropic limit of kinetic theory \citep{Rutherford:1968,Frieman:1982,Howes:2006,Schekochihin:2009}. In the derivation of the gyrokinetic equation, the crucial step is an averaging over the particle gyrophase, which leads to a theory with the following properties: the fast/whistler wave and the cyclotron resonances are discarded; all finite Larmor radius effects and collisionless dissipation via the Landau resonance are retained; and one of the dimensions of velocity in phase space is eliminated, reducing the particle distribution function from six to five dimensions. It has been previously pointed out that one cannot rule out the contribution of fast wave or whistler wave physics to solar wind turbulence, and that that therefore gyrokinetics is an incomplete description of the turbulence \citep{Matthaeus:2008a}. The novel observational analysis presented here suggests that the fast wave, in fact, does not play a statistically significant role in the turbulent dynamics of the inertial range, and that therefore a gyrokinetic approach sufficiently describes all important physical mechanisms in the solar wind inertial range. | 12 | 6 | 1206.6564 |
1206 | 1206.2640.txt | With the advent of very large volume, wide-angle photometric redshift surveys like e.g. Pan-STARRS, DES, or PAU, which aim at using the spatial distribution of galaxies as a means to constrain the equation of state parameter of dark energy, $w_{DE}$, it has become extremely important to understand the influence of redshift inaccuracies on the measurement. We have developed a new model for the anisotropic two point large-scale ($r\ga 64 h^{-1}$\,Mpc) correlation function $\xi(r_p,\pi)$, in which nonlinear structure growth and nonlinear coherent infall velocities are taken into account, and photometric redshift errors can easily be incorporated. In order to test its validity and investigate the effects of photometric redshifts, we compare our model with the correlation function computed from a suite of 50 large-volume, moderate-resolution numerical $N$-body simulation boxes, where we can perform the analysis not only in real- and redshift space, but also simulate the influence of a gaussian redshift error distribution with an absolute rms of $\sigma_z= 0.015$, $0.03$, $0.06$, and $0.12$, respectively. We conclude that for the given volume ($V_{box}=2.4 h^{-3}$\,Gpc$^3$) and number density ($\bar{n}\approx1.25\times 10^{-4} $) of objects the full shape of $\xi(r_p,\pi)$ is modeled accurately enough to use it to derive unbiased constraints on the equation of state parameter of dark energy $w_{DE}$ and the linear bias $b$, even in the presence of redshift errors of the order of $\sigma_z$ = 0.06. | The accelerated expansion of the Universe was first detected by \citet{1998AJ....116.1009R} and \citet{1999ApJ...517..565P} from the analysis of supernova type Ia observations. Over subsequent years a variety of independent data sets have confirmed this finding \citep{2005MNRAS.362..505C,2009ApJS..180..330K,2009ApJ...700.1097H,2009ApJS..185...32K,2009MNRAS.400.1643S,2010MNRAS.401.2148P,2011ApJS..192...18K}. The origin of the accelerated expansion of the Universe is one the most challenging open problems in cosmology. Several theories have been proposed to provide an explanation for this phenomenon \citep[see e.g.][for a review]{2008GReGr..40..301D}. One of the most promising solutions is the inclusion of a homogeneously and isotropically distributed fluid with a repulsive gravitational force to the energy-momentum tensor in the Einstein equations. This so-called {\it dark energy} can be characterized by its equation of state parameter $w_{\rm DE}= p_{\rm DE}/\rho_{\rm DE}$, the ratio of the pressure and energy density of this component. Current observations are consistent with the simplest solution, i.e. a cosmological constant, for which $w_{\rm DE}= -1.$ \citep[see e.g.][]{2009MNRAS.400.1643S,2010MNRAS.401.2148P,2011MNRAS.416.3017B,2011MNRAS.418.1707B,2011ApJS..192...18K,2012arXiv1202.0092M,2012arXiv1203.6616S}, although plenty of room still exists for alternative scenarios. Several observational probes have been proposed as means to obtain new clues about the nature of dark energy. Among them, the analysis of Baryonic Acoustic Oscillations (BAOs) is considered as one of the most promising alternatives \citep{arXiv:astro-ph/0609591}. BAOs are the fossil signal of the acoustic waves that propagated through the photon-baryon fluid until the epoch of recombination, when electrons and protons formed neutral hydrogen and the speed of sound dropped rapidly to zero. The maximum distance the sound waves were able to travel is called sound horizon. This scale is imprinted on the large-scale matter density fluctuations. In the two point correlation function $\xi(r)$ the signature of the BAOs shows up as a single broad bump at around $110\,h^{-1}{\rm Mpc}$ \citep{2004ApJ...615..573M}, in the power spectrum as a series of small wiggles \citep{1998ApJ...496..605,2012MNRAS.421.2656M}. Since galaxies form in the high-density peaks of the matter density field, BAOs are also present in the distribution of galaxies at later times, and have indeed been detected \citep{2005ApJ...633..560E,2005MNRAS.362..505C,2009MNRAS.396.1119C,2011MNRAS.416.3017B,2011MNRAS.418.1707B} locally and at intermediate redshift. Since they originate in the early Universe where dark energy does not play a role and their propagation is described by well understood plasma physics, the BAO signal can be used as a standard ruler. By measuring the apparent extent of the acoustic scale in the directions parallel and perpendicular to the line-of-sight it is possible to recover the redshift evolution of the Hubble parameter $H$ and the angular diameter distance $D_{\rm A}$ through a simple geometrical relation \citep{2003ApJ...594..665B}. Due to the low amplitude of the BAO signal, their analysis requires the observation of large volumes, which due to the expenditure of time spectroscopy causes, are often only accessible with photometric surveys -- which in turn comes at the price of large redshift inaccuracies. Besides this technical difficulty, which we will address in this paper, galaxy clustering differs from the linear theory predictions in a number of ways that need to be taken into account if these measurements are to be used to obtain constraints on cosmological parameters. Nonlinear growth of structure leads to coupling between Fourier modes, changing the shape of the power spectrum and correlation function. The measured clustering statistics are also affected by the gravitationally induced peculiar motions of galaxies which introduce a distortion when the distance to each galaxy is inferred from its observed redshift. Besides this, galaxies are biased tracers of the underlying dark matter density field. As a result, the correlation function of the galaxies could be a modified version of that of the mass. Besides their effect on the broad-band shape of the correlation function and power spectrum, these effects alter the shape and location of the BAO signature \citep{2008MNRAS.383..755A,2008MNRAS.390.1470S,2008PhRvD..77d3525S} and might jeopardize the success of any analysis if they are not modelled accurately. Redshift-space distortions originating from the peculiar velocities of the galaxies are unfortunately not the only distortions of the line-of-sight component of the correlation function. The measurement of redshifts are not free of errors, which add to the distortion of the clustering signal. These errors are particularly important in the case of photometric redshifts where they can be as large as $\sigma_z$ = 0.03 -- 0.04, smearing out the clustering signal along the line-of-sight, and leading to a significant reduction of its amplitude at a given scale. Examples of large photometric redshift surveys with the aim to measure the equation of state of dark energy using galaxy clustering are the Panoramic Survey Telescope And Rapid Response System Pan-STARRS \citep{2004AAS...204.9702C}, the upcoming Dark Energy Survey DES \citep{2010AAS...21547009T}, or PAU ({\bf P}hysics of the {\bf A}ccelerating {\bf U}niverse, see \citet{Benitez09}). As described above, in order to obtain constraints on both $D_{\rm A}(z)$ and $H(z)$ it is necessary to measure the full two-dimensional clustering pattern by splitting the comoving distance between the galaxies into a component perpendicular, $r_p$, and parallel, $\pi$, to the line-of-sight. Despite this, most of the theoretical analyses and observations to date have focused on angle-averaged statistics which are sensitive instead to the parameter combination $D_{\rm A}^2(z)/H(z)$. In this paper we model the full $\xi(r_p,\pi)$ by taking nonlinear structure growth, nonlinear coherent infall velocities, and redshift errors into account. The paper is structured as follows: In Section \ref{method} the analytic model and the tests performed to compare it with the results of N-body simulations are discussed in detail. The performance of the model for real-, redshift- and redshift-error space is described in Section \ref{discussion}, where it is also compared to previous work. Finally, in Section \ref{conclusions} the most important results of this paper are summarized. | In this work we developed and tested a model of the anisotropic two point correlation function $\xi(r_p,\pi)$, which we used to investigate the influence of photometric redshift errors on the measurement of the dark energy equation of state parameter $w_{{\mathrm DE}}$. We modeled $\xi(r_p,\pi)$ using third order perturbation theory \citep{1994ApJ...431..495J} to account for the nonlinear nature of the growth of structure and the nonlinear Kaiser effect \citep{2004PhRvD..70h3007S}. Redshift errors can be included in the model by convolving it with the pairwise redshift error distribution, which can easily be computed from the (known) photometric redshift errors. In order to test the validity of our model, we fit it to the mean correlation function measured from the dark matter haloes in a suite of 50 large-volume, medium-resolution $N$-body simulations (the L-BASICCS II, \citealp{2008MNRAS.387..921A,2008MNRAS.390.1470S}). Both in real and redshift space the fit yields unbiased values of the dark energy equation of state parameter $w_{{\mathrm DE}}$ and the linear bias $b$. With approximately $300\,000$ haloes per box, in real space $w_{{\mathrm DE}}$ and $b$ can be determined with an accuracy of about $12$\% and $7$\%, respectively. In redshift space these constraints become slightly weaker, $w_{{\mathrm DE}}$ can be measured with an accuracy of approximately $15$\%, and the relative error of $b$ becomes $\sim 8$\%. If only the shape information is used to infer $w_{{\mathrm DE}}$ and $b$, the errors on both will increase due to the lack of information contained in the amplitude. The relative error of the bias increases more than the relative error of $w_{{\mathrm DE}}$, since the value of the bias is mainly encoded in the amplitude (and less in the quadrupole and hexadecapole contribution to the redshift space distortion), whereas the equation of state parameter of dark energy influences both shape and amplitude likewise. In order to investigate the effect of redshift errors on the measurement, we added a small offset to one of the coordinates of the dark matter haloes, which we drew randomly from a gaussian error distribution, and convolved the model with the corresponding pairwise redshift error distribution in the direction along the line-of-sight ($\pi$). Redshift errors smear out the clustering signal and diminish its amplitude; at the same time the convolution leads to a mixing and increase of the noise of the measurement in the single pixels, because intrinsic errors are also distributed along the line-of-sight. The impact of this on the constraints on cosmological parameters is two-fold: Since the signal of the BAOs (as the main feature of the otherwise smooth correlation function) becomes weaker in the observed range of scales, its predictive power decreases -- in the case of very large redshift errors ($\sigma_z\ga 0.06$) the signal is smeared out over such a large range of distances that it completely disappears in the noise. However, since much higher accuracies can be achieved in realistic ongoing or near-future photometric surveys such as e.g. Pan-STARRS (see \citealp{Saglia2012}), this is not a cause for concern. Integrating $\xi(r_p,\pi)$ to obtain $w(r_p)$, as originally proposed by \citet{1980lssu.book.....P} as a means to overcome redshift space distortions, does not help to improve the constraints, as in real space the BAO is a {\it ring} in the $\xi(r_p,\pi)$ plane, and, when integrated, is distributed over $0.\leq r_p\la 120 h^{-1}$\,Mpc. Since it is impossible to integrate $\xi(r_p,\pi)$ to $\pi=\infty$, the resulting amplitude and shape of $w(r_p)$ depends on the choice of integration limits as well as the underlying cosmology, which adds a further complication. Secondly, the noise itself increases in the presence of redshift errors, which creates an additional difficulty. Due to the decreased signal to noise of the correlation function, the accuracy of the constraints on $w_{{\mathrm DE}}$ and $b$ decreases. In order to beat down systematics coming from cosmic variance (which is still large, even on BAO scales), it is desireable (and important) to observe the largest volumes possible at one particular redshift. Also, in order to measure a possible variation in the equation of state with lookback time, observations have to be carried out at higher redshifts as well. At this moment in time both is still only feasible with photometric redshifts. The anisotropic correlation function $\xi(r_p,\pi)$, which can be used to infer cosmological parameters like the dark energy equation of state $w_{{\mathrm DE}}$, is well suited to incorporate photometric redshifts. We have developed a model of $\xi(r_p,\pi)$ which will be able to provide {\it unbiased} constraints on $w_{{\mathrm DE}}$ and $b$ for photometric redshift surveys. The maximum redshift error for which this model will work certainly depends on the exact shape of the redshfit error distribution, the volume and number density of the survey to which it is applied. | 12 | 6 | 1206.2640 |
1206 | 1206.4088_arXiv.txt | Phase-induced amplitude apodization (PIAA) coronagraphs are a promising technology for imaging exoplanets, with the potential to detect Earth-like planets around Sun-like stars. A PIAA system nominally consists of a pair of mirrors which reshape incident light without attenuation, coupled with one or more apodizers to mitigate diffraction effects or provide additional beam-shaping to produce a desired output profile. We present a set of equations that allow apodizers to be chosen for any given pair of mirrors, or conversely mirror shapes chosen for given apodizers, to produce an arbitrary amplitude profile at the output of the system. We show how classical PIAA systems may be designed by this method, and present the design of a novel 4-mirror system with higher throughput than a standard 2-mirror system. We also discuss the limitations due to diffraction and the design steps that may be taken to mitigate them. | The challenge of imaging an Earth-like planet around a Sun-like star has two primary components: contrast and angular resolution. An Earth twin orbiting a sun twin at 1 AU at a distance of 10 parsecs will be $10^{10}$ times fainter than its host star, while being separated by only 100 milliarcseconds. For a 2m telescope with a circular aperture operating at 500nm---a reasonable expectation for a visible-band space telescope---this places the planet at $1.95 \lambda/D$ in the image plane. Here, the intensity in the point spread function (PSF) from the host star will be $\sim6\times10^{-3}$ lower than the peak of the diffraction pattern and over six orders of magnitude higher than the planet, making detection of the planet virtually impossible. To overcome this, a number of optical systems---including coronagraphs, occulters, and interferometers---have been proposed to reduce the residual stellar light in the regions of the image plane where we wish to look for planets. One of the most promising coronagraphs for space applications is the phase-induced amplitude apodization (PIAA) coronagraph \mycitep{Guy03, Tra03, Guy05, Van05}, which uses a pair of mirrors to losslessly remap incident light to apodize the beam while retaining high throughput. In addition, incident wavefronts with an angle to the optical axis of the system are displaced disproportionately from the optical axis, providing an effective magnification of the inner working angle. A number of these systems have been designed \mycitep{Plu06} and tested \mycitep{Bel09, Ker09, Guy10, Bal10, Pue11a}, in conjunction with apodizing elements to reduce manufacturing requirements on mirrors or mitigate diffraction from mirror edges. As shown in \mycitet{Van06}, without at minimum an apodizer before the mirrors (a \emph{pre-apodizer}), the effects of diffraction from the edges of the mirrors limit the achievable contrast severely, to $10^{-5}$. Apodizers after the mirrors (a \emph{post-apodizer}) have also proved useful for performing a portion of the beam-shaping, at the expense of some throughput; using a post-apodizer can reduce the manufacturing requirements on the mirrors, particularly the curvatures required on the edges of some designs \mycitep{Plu06}. While the tools for accurately modeling propagation through PIAA systems have become fairly sophisticated and robust \mycitep{Bel06a, Pue09, Kri10, Pue11}, the tools for finding new apodizations and mirror shapes are not as developed; in particular, the interaction between multiple remapping and apodizing components in the design has not been fully explored. In this paper, we discuss the propagation of radially-symmetric amplitude functions through PIAA systems, in particular setting down a series of equations in Sec. \ref{sec:3f} which allow for the arbitrary specification of amplitude profiles throughout the system. This permits mirrors and apodizers to be designed explicitly to work jointly as a system; examples of this are shown in Sec. \ref{sec:app}. Applications to design and modeling of four-mirror systems are discussed in Sec. \ref{sec:4m}, including their interactions with apodizers in the system. | In this work, we have presented three sets of equations which allow apodizations and mirror shapes to be designed exactly given choices for the remaining elements. We expect in the near future to provide explicit representations for the residual diffraction in terms of the amplitude profiles presented here, and the derivation of these equations from an approximation to the diffraction integral provides a solid foundation for this future work. (This approximation was also shown to reproduce the geometric-optics amplitude distribution for a uniform input field.) A number of examples demonstrating the use of these equations have been provided; it is our hope that these will prove to be useful tools for subsequent PIAA designs. We have also used these tools to design and analyze systems with multiple pairs of PIAA mirrors, to allow significantly-improved throughput as well as smaller inner working angle. We expect to verify the performance of this class of system experimentally in the upcoming months; initial progress may be found in \mycitet{Bal11}. | 12 | 6 | 1206.4088 |
1206 | 1206.1452_arXiv.txt | We present the first orbit--integrated self force effects on the gravitational waveform for an I(E)MRI source. We consider the quasi--circular motion of a particle in the spacetime of a Schwarzschild black hole and study the dependence of the dephasing of the corresponding gravitational waveforms due to ignoring the conservative piece of the self force. We calculate the cumulative dephasing of the waveforms and their overlap integral, and discuss the importance of the conservative piece of the self force in detection and parameter estimation. For long templates the inclusion of the conservative piece is crucial for gravitational--wave astronomy, yet may be ignored for short templates with little effect on detection rate. We then discuss the effect of the mass ratio and the start point of the motion on the dephasing. | The detection of gravitational waves (henceforth GW) and the onset of the new field of gravitational-wave astronomy is one of the most exciting challenges for science in the XXI century, completing what is sometimes alluded to as Einstein's Unfinished Symphony. The detection of GW will open a new window onto the universe, that in addition to revealing exciting information on exotic systems such as black holes or cosmic strings is expected to also unravel as yet unexpected sources. % One of the interesting sources for low--frequency GW are the so called I(E)MRI sources, or Intermediate (Extreme) Mass Ratio Inspirals. Those are the GW emitted by a system including a smaller compact object whose orbit decays into a much larger massive black hole (MBH). Typical sources are stellar mass black holes inspiraling into a supermassive black hole, like those residing at the center of galaxies, and also IMBHs (intermediate mass black holes) inspiraling into MBHs. The importance of such sources is that because of the extreme mass ratio the smaller compact object can be viewed as a test particle, thus probing the spacetime of the larger black hole and its surroundings. {\it Inter alia}, such sources will allow us to test directly the Kerr hypothesis, and allow us to map the spacetime surrounding such exotic objects. Moreover, the detection of I(E)MRIs will allow us to determine the mechanisms that shape stellar dynamics in galactic nuclei with unprecedented precision \cite{Amaro-Seoane2012}. The orbits of I(E)MRIs are typically highly relativistic, and exhibit exciting phenomena, e.g.~extreme periastron and orbital plane precessions. Because the orbital evolution time scale (``radiation reaction time scale'') is much longer than the orbital period(s), over short time scales the orbit is approximately geodesic, yet on long time scales it deviates strongly from geodesic motion of the background. Instead, the smaller objects moves along a geodesic of a perturbed spacetime. Alternatively, one may construe the orbit as an accelerated, non-geodesic motion in the spacetime of the unperturbed central object, where the acceleration is caused by the self force (henceforth SF) of the smaller object \cite{Poisson-Pound-Vega}. Detection and parameter estimation of GW from I(E)MRIs relies on the construction of theoretical templates. A number of approximation schemes for such templates are available. Firstly, the energy balance approach (``the radiative approximation'') uses balance arguments for otherwise conserved quantities, and relates the flux in these quantities to infinity and down the event horizon of the black hole with the particle's orbit, so that the latter can be adjusted to agree with the fluxes \cite{Hughes}. % As the orbital evolution time scale is typically much longer than the orbital period(s), the radiative approximation is very satisfactory during the adiabatic phase of the motion. As the particle's orbit is affected by the fluxes away from it, when the orbit is not stationary one encounters complex retardation effects. % Most currently--available EMRI waveforms have been obtained by such an approach. This approach, however, ignores conservative effects that do not register in the constants of motion. These retardation effects are completely avoided when one considers a local approach to orbital evolution in terms of the SF. (One should bear in mind, however, that the SF itself is a non-local quantity, with contributions arising from the quasilocal neighborhood of the particle and possibly beyond \cite{Anderson}.) In addition, the local approach to the calculation of orbital evolution via the SF is not restricted to the adiabatic regime, it avoids the complications associated with the rate of change of the Carter constant, and, most importantly, it includes also conservative effects that are discarded when one uses balance arguments. Over the last decade much progress has been made in the computation and understanding of the SF in General Relativity (for recent reviews of the self force in General Relativity see \cite{Poisson-Pound-Vega,barack:2009}). The computation of the fully relativistic SF allows one to include conservative effects in the waveform templates, and study the importance of the conservative effects. True self consistent orbit and waveforms include the instantaneous solution of the coupled SF integrated equations of motion and the perturbation equations, or equivalently the interaction of the particle with its own field over its half--infinite past world line \cite{gralla-wald}. Very recently, for the scalar field toy model, such self consistent Schwarzschild orbits and waveforms were presented \cite{diener}. Here, we are making the simplifying assumption that the effects of the difference between the SF that is calculated for the actual orbit (the self consistent approach \cite{Pound,gralla-wald}) and that which is calculated for a geodesic of the same instantaneous orbital parameters, is smaller than the effects of the latter and hence negligible at first order. This approximation is valid for as long as the orbital evolution is adiabatic, that is as long as the orbital evolution time scale is much longer than the orbital period(s). In a Schwarzschild background of mass $M$, the adiabatic approximation holds when the mass ratio $\eta:=\mu/M$ is such that $\varepsilon\gg \eta^{1/2}$, where $\varepsilon$ measures the distance to the innermost stable orbit, specifically $\varepsilon=p-6-2\epsilon$ where $p$ is the semilatus rectum and $\epsilon$ is the orbital eccentricity \cite{cutler}. In practice, our approximation is to a leading order in $\eta$ beyond geodesic motion. We neglect terms that are linear in second--order SFs, although our method is amenable to their inclusion when they become available. This approximation is valid for at least a part of the relevant parameter space \cite{burko04}, but as their inclusion would contribute linearly to the dephasing, the contribution of the conservative piece of the SF (hereafter CSF) may be isolated as is done here. Using true self consistent waveforms will both produce more accurate waveforms, and allow us to test the accuracy of this approximation. Most importantly, our approach allows us to see for the first time the effect of the CSF on GW emitted from IMRI sources. We present here the first waveforms obtained with inclusion of the CSF, and study its effect within the simple class of quasi--circular orbits around a Schwarzschild black hole. Specifically, we study the effect of the system's mass ratio on the dephasing that occurs when one neglects the CSF. We find weak dependence of the dephasing on the mass ratio, in accord with expectations based on the scaling of the number of orbits with the inverse of the mass ratio, and the scaling of the dephasing effect of the CSF per orbit with the mass ratio. We also find that the dephasing depends quadratically on the initial point of the motion for the range of parameters we tested. We reiterate that second--order dissipative effects are ignored in this Paper. Their inclusion will guarantee the full consistency of the model, and will be comparable to the self-consistent approach. The inclusion of the second-order dissipative effects awaits further development to both theory and computational techniques. The organization of this Paper is as follows: In Section \ref{method} we discuss the computational and numerical methods that we use. In Section \ref{results} we discuss our results for the orbits (\ref{orbit}), the waveforms (\ref{waveform}) and the dependence of the dephasing on the mass ratio and the initial point of the motion (\ref{varying}). | 12 | 6 | 1206.1452 |
|
1206 | 1206.6334_arXiv.txt | {Direct imaging of circumstellar disks at high angular resolution is mandatory to provide morphological information that bring constraints on their properties, in particular the spatial distribution of dust. This challenging objective was for a long time in most cases, only within the realm of space telescopes from the visible to the infrared. New techniques combining observing strategy and data processing now allow very high contrast imaging with 8-m class ground-based telescopes ($10^{-4}$ to $10^{-5}$ at $\sim$1") and complement space telescopes while improving angular resolution at near infrared wavelengths. } {We carried out a program at the VLT with NACO to image known debris disks with higher angular resolution in the near IR than ever before in order to study morphological properties and ultimately to detect signpost of planets.} {The observing method makes use of advanced techniques: Adaptive Optics, Coronagraphy and Differential Imaging, a combination designed to directly image exoplanets with the upcoming generation of "planet finders" like GPI (Gemini Planet Imager) and SPHERE (Spectro-Polarimetric High contrast Exoplanet REsearch). Applied to extended objects like circumstellar disks, the method is still successful but produces significant biases in terms of photometry and morphology. We developed a new model-matching procedure to correct for these biases and hence to bring constraints on the morphology of debris disks.} {From our program, we present new images of the disk around the star HD\,32297 obtained in the H (1.6$\mu m$) and Ks (2.2$\mu m$) bands with an unprecedented angular resolution ($\sim$65 mas). The images show an inclined thin disk detected at separations larger than $0.5-0.6"$. The modeling stage confirms a very high inclination ($i=88\degb$) and the presence of an inner cavity inside $r_0\approx 110$\,AU. We also found that the spine (line of maximum intensity along the midplane) of the disk is curved and we attributed this feature to a large anisotropic scattering factor ($g \approx 0.5$, valid for an non-edge on disk).} {Our modeling procedure is relevant to interpret images of circumstellar disks observed with Angular Differential Imaging. It allows both to reduce the biases and to provide an estimation of disk parameters. } | \begin{table*}[t] \caption{Log of observations summarizing the parameters: the filter used, the time of observation, the integration time of a single frame (DIT), the number of frames per data cube (NDIT), the number of cubes (Nexp), the seeing variation, the mean correlation time $\tau_0$, the amplitude of the parallactic angle variation and the observation time on target.} \begin{center} \begin{tabular}{ccccccccccc} \hline\hline Date & Filter & UT start/end & DIT (s) & NDIT & Nexp & seeing (") & $\tau_0$ (ms) & Parall. amplitude ($^\circ$) & { on target (s)} \\ \hline 2010.09.25 & H & 07:23:40 / 08:45:15 & 2 / 3 & 15 / 10 & { 68/13} & 0.70 - 1.28 & 1.7 & 24.81 & { 1710} \\%-137.37 / -162.18 \\ 2010.09.25 & Ks & 09:04:08 / 10:00:57 & 2 & 15 & 57 & 0.59 - 0.88 & 2.0 & 24.18 & { 2430} \\ \hline % \end{tabular} \end{center} \label{tab:log} \end{table*} Many stars have photometric IR excess attributed to the presence of circumstellar dust. Among the circumstellar disk evolution, the debris disk phase is the period where most of the primordial gas { has been} dissipated or accreted onto already formed giant planets. The dust content responsible for the IR excesses is thought to be the result of regular mutual collisions between asteroid-like objects and/or evaporation of cometary bodies \citep{Wyatt2008, Krivov2010}. This activity, presumably triggered by unseen planets, is continuously replenishing the system with observable dust grains that, in turn, trace out the presence of a planetary system. The recent detection by direct imaging of planets in circumstellar disks around main sequence A-type stars like $\beta$ Pictoris \citep{Lagrange2010} and HR 8799 \citep{Marois2008} or the undefined point source around Fomalhaut \citep{Kalas2008, Janson2012} has clearly revealed the suspected intimate relation between { planet} formation and debris disks structures \citep{Ozernoy2000, Wyatt2003, Reche2008}. Moreover, the most prominent structures in these { specific} disks (warp, annulus, gap, offset) can be accounted for by the mass and orbital properties of { these} associated planets. { Similar features are seen in other debris disks where planets are not yet imaged.} Many disks are thought to have internal cavities attributed to dust sweeping by planets. From a face-on to an edge-on geometry these structures are more or less easily identified. Therefore, a morphological analysis is mandatory to study disk evolution, planetary formation and { the likely} presence of planets. HD\,32297 is a young A5 star identified with an IR excess ($L_{IR}/L_*=0.0027$) by IRAS and later resolved as a { inclined} circumstellar debris disk by \citet{Schneider2005} in the near IR with NICMOS on HST. Observations in the visible (R band) were obtained by \citet{Kalas2005} and reveal a circumstellar nebulosity at distances larger than 560\,AU indicating a broad scattered-light region extending to 1680\,AU. The brightness asymmetries and a distorted morphology suggests the interaction with the interstellar medium. HST data obtained at 1.1$\mu$\,m probe the inner region below $\sim$400\,AU and as close as 34\,AU from the star (corresponding to respectively 3.3" and 0.3"). \citet{Schneider2005} derived a Position Angle ($PA$) of $47.6\pm1.0$\deg and an inclination of about 10\deg from edge-on. A surface brightness asymmetry is also observed in the South West (SW) side of the disk which is about twice brighter than the North East (NE) side at a projected angular separation of about 0.3-0.4". Observations at mid-IR wavelengths (11.7 and 18.3$\mu$\,m) by \citet{Moerchen2007} indicate a dust depletion closer than $\sim$70\,AU suggestive of a inner cavity like the one found in HR\,4796 \citep{Schneider1999} or HD\,61005 \citep{Buenzli2010} but with a less steep decrease of the density. However, this radius is not well constrained in the mid IR as the equivalent resolution corresponds to 35 to 50\,AU at the HD\,32297 distance. Additional colors were obtained by \citet{Debes2009} at 1.60 and 2.05$\mu$\,m again with NICMOS which confirms the disk orientation ($PA=46.3\pm1.3$\deg at NE and $49.0\pm2.0$\deg at SW). The surface brightness profiles confirm the asymmetries and reveal a slope break at 90\,AU (NE) and 110\,AU (SW). If { similar to some} other disks like $\beta$ Pictoris \citep{Augereau2001} and AU Mic \citep{Augereau2006} this break indicates the boundary of a planetesimals disk, it is marginally compatible with the value (70\,AU) derived by \citet{Moerchen2007} but the measurements may not be directly comparable to that of \citet{Schneider2005}. Also, it is shown that the disk midplane is curved at separations larger than 1.5" (170\,AU), possibly a result of the interaction with the interstellar medium. This probably impacts the $PA$ measurement. More recently, \citet{Mawet2009} achieved very high contrast imaging of the disk with a small (1.6\,m) well corrected sub-aperture at the Palomar in Ks and confirmed the colors and brightness asymmetry, but the lower angular resolution did not allow a detailed morphological study. At a distance of $113 \pm 12$\,pc \citep{Perryman1997}, the disk of HD\,32297 is of particular interest to probe the regions where planets would have formed in the 10-AU scale providing high angular resolution and high contrast are achieved simultaneously. In this paper, we present new diffraction limited images of the HD\,32297 debris disk in the near IR (1.6 and 2.2 $\mu$m) using an 8-m class telescope { with adaptive optics}, therefore with a significant improvement with respect to former observations (a factor of $\sim$3 { for the angular resolution} at the same wavelengths). We note that similar observations obtained at the Keck telescope in the same near IR bands are presented by \citet{Esposito2012}. Section \ref{sec:observ} presents the results of the observation. In section \ref{sec:morpho} we analyze the reliability of a particular structure observed in the image. Section \ref{sec:photom} provides the Surface Brightness profiles of the disk before biases are taken into account. The core of this paper is presented in the section \ref{sec:model} where we attempt to identify the parameters of the disk. In section \ref{sec:photomcorrec} the Surface Brightness corrected for the observational bias is derived. Finally, we present constraints on the presence of point sources within the disk. | We { present} high contrast images of the debris disk around HD\,32297 obtained with NACO at the VLT in the near IR H and Ks bands. To improve the attenuation of the starlight and the speckle rejection we combined phase mask coronagraphy and ADI { algorithms}. The disk { nearly edge-on} is detected in the two filters { at radial distances} from 0.65" to 2" (74 to 226 \,AU at the distance of 113\,pc) owing to various ADI algorithms. ADI is a very efficient method to achieve detection but induces some { photometric} biases that have to be { corrected for in the data analysis for a proper interpretation}. We observed a deviation from the midplane near 0.8" in the NE side that is clearly not an ADI bias. We focused our effort on the morphological modeling of the disk using a numerical code GRaTer, which produces synthetic scattered light images of debris disks. Following previous observations, we considered a disk containing a planetesimal belt with dust distribution inward and outward. We generated a 5-parameter grid of models and applied several sorts of minimizations to constrain the disk morphology. As a first result, we found { in Ks and H bands} a higher inclination than previously published for $1.1\mu m$ (88$\degb$ instead of 80$\degb$). In our images, the vertical extent of the disk together with the disk parameters found from modeling are definitely not consistent with such a disk inclination. However, we suspect that the isophotal ellipse fitting used in former data \citep{Schneider2005} can be biased by the strong interaction with the interstellar medium as revealed by \citet{Debes2009}, especially when considering the lower angular resolution and conversely the higher sensitivity to faint structures of NICMOS/HST as opposed to NACO. In addition, we did not found any significant NE/SW brightness asymmetry like in \citet{Schneider2005} and \citet{Mawet2009} although we are { also observing at near IR} wavelengths { too}. Furthermore, the minimization of the disk parameters allowed to constrain the radius of the planetesimal belt at 110\,AU. A NE/SW asymmetry is suspected, possibly due to an offset of the planetesimal belt with respect to the star. We showed that the deviation from the midplane is { likely} the result of anisotropic scattering. One edge of the disk is therefore brighter. Whether it is linked to the properties of the dust or to the geometry (the hypothetical offset of the belt) and hence to a gravitational perturber is yet undetermined. It is remarkable that among the known debris disks, three of them surrounding A-type stars, namely, $\beta$ Pic, Fomalhaut, and HD\,32297 share the same position of the planetesimal belt near 100\,AU as a possible indication of a common architecture. Intensity related parameters like the brightness slopes of the inner and outer parts are more difficult to constrain and the values we determined should be taken with caution as the uncertainty can be large. However, it gives an order of magnitude indicating that the outer parts have a power law index close to that expected for dust grains expelled from the system by radiation pressure. As for the inner parts, there is a significant difference between Ks and H band filters, but overall, the steepness is smaller than some other ring-like debris disks as already noted by \citet{Moerchen2007}. Armed with a detailed morphological characterization it is now possible to start a dynamical study to understand the properties of the debris disk around HD\,32297 in regards of planetary formation or gravitational perturbation as it was performed in other cases \citep{Wyatt2005, Chiang2009}. Finally, the observations of debris disks from ground-based telescopes { in the near IR} are achieving comparable or even better angular resolution than from HST, owing to { adaptive optics. A similar statement can be made about contrast} owing to a combination of high contrast imaging techniques (coronagraphy, ADI). The next generation of extreme adaptive optics instruments, the so-called planet finders \citep{Beuzit2008, Macintosh2008, Hodapp2008}, will not only revolutionize the search and characterization of young long-period planets but also the field of debris disk science pending the arrival of the James Webb Space Telescope \citep{Gardner2006}. | 12 | 6 | 1206.6334 |
1206 | 1206.6102_arXiv.txt | We have performed 323 MHz observations with the Giant Metrewave Radio Telescope of the most promising candidates selected from the MACS catalog. The aim of the work is to extend our knowledge of the radio halo and relic populations to $z>0.3$, the epoch in which massive clusters formed. In MACSJ1149.5+2223 and MACSJ1752.1+4440, we discovered two double-relic systems with a radio halo, and in MACSJ0553.4-3342 we found a radio halo. Archival Very Large Array observations and Westerbork Synthesis Radio Telescope observations have been used to study the polarization and spectral index properties. The radio halo in MACSJ1149.5+2223 has the steepest spectrum ever found so far in these objects ($\alpha \geq$ 2). The double relics in MACSJ1149.5+2223 are peculiar in their position that is misaligned with the main merger axis. The relics are polarized up to 30\% and 40\% in MACSJ1149.5+2223 and MACSJ1752.040+44, respectively. In both cases, the magnetic field is roughly aligned with the relics' main axes. The spectra in the relics in MACSJ1752.040+44 steepen towards the cluster centre, in agreement with model expectations. X-ray data on MACSJ0553.4-3342 suggests that this cluster is undergoing a major merger, with the merger axis close to the plane of the sky. The cores of the disrupted clusters have just passed each other, but no radio relic is detected in this system. If turbulence is responsible for the radio emission, we argue that it must develop before the core passage. A comparison of double relic plus halo system with cosmological simulations allows a simultaneous estimate of the acceleration efficiencies at shocks (to produce relics) and of turbulence (to produce the halo). | A fraction of galaxy clusters host diffuse radio emission, that is not connected to any of the cluster radiogalaxies. These radio sources are classified as radio halos and radio relics, depending on their location and morphology. In all these sources relativistic particles need to be (re)accelerated in order to produce the observed radio emission, even though the underlying physical mechanisms are likely to be different. Radio halos permeate the central Mpc of galaxy clusters, and in some cases the radio emission roughly follows the X-ray emission from the thermal gas \citep[see review by][and references therein]{2008SSRv..134...93F}. Radio halos are characterized by a steep radio spectrum,\footnote{We define here the spectrum as $S(\nu) \propto \nu^{-\alpha}$. } with $\alpha\geq 1.2$, and their power at 1.4 GHz correlates with the X-ray luminosity of the host cluster \citep{2000ApJ...544..686L,Giovannini09}. The origin of radio halos is still unknown, although a clear connection between the presence of a radio halo and the merging state of the host cluster is present \citep{2001ApJ...553L..15B,2010ApJ...721L..82C,2011ApJ...740L..28B,2012MNRAS.421L.112B}. The models proposed so far can be divided into two classes: ``hadronic models'' \citep[e.g][]{1980ApJ...239L..93D,2010ApJ...722..737K,2011A&A...527A..99E} and ``re-acceleration models''. It was recently shown that hadronic models fail to reproduce the observed radio emission of the Coma cluster, once the upper limits by FERMI and the magnetic field estimate from Faraday Rotation Measures are combined \citep{2010MNRAS.401...47D}. Recent LOFAR radio observations of the cluster Abell 2256 have shown that the spectral flattening predicted at low frequencies by re-acceleration models are not observed \citep{2012arXiv1205.4730V}. Hence, the question about the origin of radio halos is still open, and it is likely that more complex scenarios have to be considered.\\ Radio relics are irregularly shaped radio sources, located at the outskirts of galaxy clusters. They usually have an arc-like structure, and are found to be polarized at 10 - 80\% level \citep[see review by][and references therein]{2011SSRv..tmp..138B}. Like radio halos, radio relics are characterized by a steep spectrum and a weak surface brightness at 20 cm that makes their detection difficult. Their origin is still unclear, but there is a common consensus that they are related to merger shocks. A merging shock could accelerate particles via Diffusive Shock Acceleration (DSA), and amplify the magnetic field strength in the shock region, hence producing radio synchrotron emission \citep[e.g.][]{1998A&A...332..395E,2012arXiv1204.2455I}. Some clusters show a relic and a radio halo, like the Coma cluster \citep{1993ApJ...406..399G}, others, such as Abell 115 only host a single relic \citep{2001A&A...376..803G}, a few objects have been found where double relics are present \citep[e.g. Abell 1240 and Abell 2345][]{2009A&A...494..429B}. Finally, in a couple of objects, both double relics and a radio halo have been found \citep[e.g. CIZA J2242.8+5301][]{2010Sci...330..347V}. The combinations with which halos and relics are found, and the number of these objects itself, are likely to be affected by detection limits of current instruments. In the last years, efforts have been made from both observational and theoretical sides to better understand the properties of radio relics. From a theoretical point of view, the properties of merging shocks, and the acceleration mechanism required to explain the radio emission, have been investigated in several works \citep{pf06,2007MNRAS.375...77H,2009MNRAS.393.1073B,2009A&A...504...33V,va10kp,2011ApJ...735...96S,2011ApJ...734...18K,2012MNRAS.420.2006N,Vazza12,2012arXiv1204.2455I}. Recently, \citet{2012MNRAS.420.2006N} have used hydrodynamic simulations to predict the number of observable relics as a function of the redshift. They find that either more relics with $z>0.3$ should be found than known to date (only a couple have been observed so far), or that the ratio $B/n_e$ of the magnetic field $B$ over the thermal gas density, $n_e$, changes with $z$. The ``lack'' of relics at $z>0.3$ can be however due to the lack of deep radio observations of distant galaxy clusters. Investigating the properties of galaxy clusters in this redshift range would hence provide information on the physics of the ICM itself. \\ So far, the properties of radio halos and relics have been mainly studied in galaxy clusters at low redshift ($z<0.3$). \citet{Giovannini09} have analyzed the radio emission from a sample of galaxy clusters at $z<0.2$, finding that at the NVSS detection limit, the percentage of galaxy clusters with diffuse sources is in the range 6-9\% in clusters with X-ray luminosity $L_x>10^{44}$ erg s$^{-1}$, and grows to $\sim$40\% when only clusters with $L_x>10^{45}$ erg s$^{-1}$ are considered. \citet{2008A&A...484..327V} have searched for the presence of radio halos in a complete sample of X-ray luminous galaxy clusters, finding that the fraction of clusters with a radio halo is $\sim$30\% in the redshift range 0.2 - 0.4 (with most of the objects having $z \leq 0.3$) increasing up to $\sim$40\% when clusters with $L_x>8 \cdot 10^{44}$ erg s$^{-1}$ are considered.\\ The aim of the present work is to extend the sample of diffuse radio sources to $z > 0.3$, the epoch when massive clusters formed. X-ray selected samples compiled from ROSAT data, such as REFLEX \citep{2004A&A...425..367B} and eBCS \citep{2000MNRAS.318..333E} have probed the radio properties of galaxy clusters in the local Universe \citep[e.g.][]{2008A&A...484..327V,Giovannini09}. At higher redshifts the MACS project has compiled the first large X-ray selected sample of clusters that are both X-ray luminous and distant \citep{Ebeling01}. It consists of 128 clusters at $z>0.3$. The MACS catalogs of \citet{Ebeling07,Ebeling10} comprise all clusters at $z>0.3$ with a nominal X-ray flux larger than $2 \, \cdot 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (0.1--2.4 keV) in the ROSAT Bright sources catalog \citep{Voges99}, and 12 clusters at $z>0.5$ down to an X-ray flux limit of $1\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$. A third MACS subsample was recently published by \citet{2012MNRAS.420.2120M}. We have inspected the archival Chandra/XMM-Newton data of the clusters contained in the MACS catalogs, and we have selected those that show a disturbed X-ray morphology. To start with, we have observed four promising candidates that show hints of diffuse emission from the NVSS \citep{NVSS}. However, there are further candidates for diffuse emission in the MACS sample that will be followed up in future work. \\ The paper is organized as follows: The observations and the data reduction are described in Sec. \ref{sec:Obs}, followed by the results in Sec. \ref{sec:results}. We discuss our results in Sec. \ref{sec:discussion}, and compare our observations with mock radio images obtained with cosmological simulations. In Sec. \ref{sec:correlations}, we present new correlations for the radio halos and double-relic systems. Finally, we conclude in Sec. \ref{sec:conclusions}.\\ Throughout this paper we assume a concordance cosmological model $\rm{\Lambda CDM}$, with $H_0=$ 72 km s$^{-1}$ Mpc$^{-1}$, $\Omega_M=$ 0.27, and $\Omega_{\Lambda}=$ 0.73. \begin{table*} \caption{Observation details} % \label{tab:radioobs} \centering \begin{tabular}{|c c c c c|} % \hline\hline Cluster & RA & DEC & $z$ & Observation Date \\ & (J2000) & (J2000)& & \\ \hline MACSJ1149.5+2223 & 11 49 34.3 & 22 23 42 & 0.544$^{c}$ & 15-AUG-201 \\ MACSJ1752.0+4440 & 17 52 01.5 & 44 40 46 & 0.366$^{b}$ & 27-JUN-2011 \\ MACSJ0553.4$-$3342 & 05 53 26.7 & -33 42 37 & 0.431$^{d}$ & 15-AUG-2010 \\ MACSJ1731.6+2252 & 17 31 40.1 & 22 52 39 & 0.389$^{a}$ & 26-JUN-2011 \\ \hline \hline \multicolumn{5}{l}{\scriptsize Col. 1: cluster name; Col. 2 and 3: Right ascension and Declination (J2000) from RASS.}\\ \multicolumn{5}{l}{\scriptsize Col. 4: redshift ($a:$ \citet{Ebeling10}, $b:$ \citet{2003MNRAS.339..913E}, $c:$\citet{Ebeling07}, $d:$ \citet{2012MNRAS.420.2120M})}\\ \multicolumn{5}{l}{\scriptsize Col. 5: Observation date.} \end{tabular} \end{table*} | \label{sec:conclusions} We have presented new radio observations of massive galaxy clusters located at high redshift ($z>0.3$). Our GMRT observations combined with WSRT and archival VLA observations have lead to the discovery of two radio halos, 2 double-relic systems and a candidate radio halo. Our results can be summarized as follows: \begin{itemize} \item{Two radio relics and a candidate radio halo are detected in the cluster MACSJ1149.5+2223, at $z=0.54$. This is the most distant cluster with double relics detected so far, and it shows peculiar features that have never been observed in other galaxy clusters. The relics have a total extent of $\sim$ 800 kpc and are located at 1.1 and 1.4 Mpc from the cluster centre, where the X-ray brightness drops down. The position of the relics is peculiar, and misaligned with the cluster merger axis. We argue that this peculiar feature could be a sign of additional merging in this complex system. In the framework of DSA, the integrated spectral indices of the relics indicate Mach numbers $M\sim 3$ and 4.6. The relics show an average polarization of 5\%, with polarization values up to 30\%. The polarization vectors suggest a magnetic field roughly oriented along the relics' main axes. The candidate radio halo has a total extent of $\sim$1.3 Mpc, and a radio flux density of $\sim$29 mJy at 323 MHz. The radio halo is barely visible in the VLA image. The spectral index $\alpha$ is larger than 2, making it the radio halo with the steepest spectrum known so far.} \item{In the cluster MACSJ1752.0+4440 ($z=$ 0.366) two radio relics and a radio halo are detected at 323 MHz. Their existence was already suggested by \citet{2003MNRAS.339..913E} and confirmed by WSRT observations \citep{2012arXiv1206.2294V}. The relics have a total extent of 1.1 and 0.9 Mpc, and are symmetric about the cluster centre. The relics in this cluster are a text-book example for shock acceleration. The spectral index trends derived between 1.7 GHz and 323 MHz show a clear steepening toward the cluster centre, possibly tracing the particle aging as we move far from the shock front. Assuming that the flattest spectral indices are representative of the injected spectrum of the particles, we derive $M \sim$ 4.6 and 2.8, in agreement with Mach numbers expected for merging shocks in the outskirts of galaxy clusters. The relics are polarized on average at a 20\% level, and show stronger polarization - up to 40\% - in the outer parts. The polarization vectors indicate an ordered magnetic field mainly aligned with the relic main axes, in line with model predictions. The radio halo has an irregular morphology. Its emission is elongated along the NE-SW direction, connected with the relics' emission. The brightest radio halo regions do not correspond to the brightest X-ray regions of the cluster. The spectral index of the radio halo between 323 MHz and 1.7 GHz is $\alpha=1.33 \pm$ 0.07.} \item{A radio halo is discovered in the cluster MACSJ0553.4-3342, at $z=0.431$. Its total extent is $\sim$ 1.3 Mpc. The halo is elongated in the EW direction, along the merger axes. The merger between two sub-clusters is likely to lie close to the plane of the sky, and the two cores have just passed each other. Given the small separation between the two X-ray cores, we derive that they passed each other $\sim$ 0.08 Gyr ago. If turbulent (re)acceleration is responsible for the halo emission, then turbulent motions must have developed well in time before the impact of the two cores. While the merger is violent enough to power a radio halo, we note the absence of radio relics in this systems and speculate about the reasons. } \item{From a set of high-resolution cosmological simulations we have picked a cluster whose projected X-ray surface brightness and morphology is similar to MACSJ1752.0+4440. We have assumed that the energy is converted from shocks to radio emission with an efficiency $A_{\rm S}$, defined such that the radio power, $P_{\rm relic}$ is equal to $A_{\rm S} \cdot \Phi_{\rm shock}$, and for turbulence with an efficiency $A_{\rm t}$ such that the halo power is $P_{\rm halo} \sim A_{\rm t} \Phi_{\rm turb}$, where $\Phi_{\rm shock}$ and $\Phi_{\rm turb}$ is the energy flux from relativistic electrons accelerated by shocks and turbulence, respectively. Adopting simple recipes, we are able to reproduce both the radio power and the observed morphology of the relic emission by assuming $A_s \approx 10^{-5}$. In order to match the power and morphology of the radio halo we require instead $A_s \approx 0.05$. This conversion factor corresponds to an efficiency of turbulent acceleration of $\eta \sim 10^{-4}$. In both cases only a very tiny fraction of the energy supplied by the merger is required to explain the observed radio emission.} \item{We have investigated the LLS-$\alpha$, distance-$\alpha$, $L_{1.4 \rm{GHz}}$- distance, $LLS-$distance, and $L_{1.4 \rm{GHz}}-z$ correlations for radio relics. We have considered only double-relics systems to minimize projection effects. We find no statistical evidence for such correlation at 5\% level of significance (2-tail test). We have found a correlation between double relics LLS and the relic radio power (LLS-$L_{1.4 \rm{GHz}}$ correlation). We note that the correlations $L_{1.4 \rm{GHz}}$- distance, $LLS-$distance, and $L_{1.4 \rm{GHz}}-z$ would be accepted if we applied a 1-tail test at 5\% level of significance. We do not consider here such test to be accurate enough to probe a robust statistical correlation, but we argue that complete samples of double-relic clusters are needed to understand whether these correlations are present or not.} \item {Collecting published data on radio halos, we find that the power of radio halos correlates with the redshift of the host cluster ($L_{1.4 \rm{GHz}}-z$ correlation), Since the most powerful mergers are expected at $z>0.3-0.4$, this correlation indicates a link between the energy injected in the ICM and the particle acceleration efficiency and/or magnetic field amplification. The LLS of the radio halo is also found to correlate with the cluster's redshift (LLS-$z$ correlation). This correlation is not easily explained since ultra-relativistic particles lose more energy through the IC mechanism as the redshift increases. Although several observational biases can affect such correlations, we argue that, if confirmed, they would provide some clues about the redshift-dependence of the particle acceleration mechanism.} \end{itemize} \bigskip {\bf Acknowledgments} We thank the referee, Shea Brown, for his useful comments on the manuscript. AB, MB, FV, MH, and UK acknowledge support by the research group FOR 1254 funded by the Deutsche Forschungsgemeinschaft: ``Magnetization of interstellar and intergalactic media: the prospect of low frequency radio observations''. We thank the staff of the GMRT that made these observations possible. GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. This research has made use of the NASA/IPAC Extragalactic Data Base (NED) which is operated by the JPL, California institute of Technology, under contract with the National Aeronautics and Space Administration. We thank G. Macario, G. Brunetti, and R. Cassano for useful discussions. FV acknowledges the usage of computational time under the INAF-CINECA agreement, and acknowledges C. Gheller for fruitful collaboration in the production of the simulations. HE acknowledges financial support from SAO grants GO1-12153X. | 12 | 6 | 1206.6102 |
1206 | 1206.2451_arXiv.txt | {We report on the discovery of an infrared emission band in the Spitzer spectrum of the S-type AGB star NP Aurigae that is caused by TiO molecules in the circumstellar environment. We modelled the observed emission to derive the temperature of the TiO molecules ($\approx 600$\,K), an upper limit on the column density ($\approx$ 10$^{17.25}$\,cm$^{-2}$) and a lower limit on the spatial extent of the layer that contains these molecules. ($\approx$\,4.6\,R$_{\star}$). This is the first time that this TiO emission band is observed. A search for similar emission features in the sample of S-type stars yielded two additional candidates. However, owing to the additional dust emission, the identification is less stringent. By comparing the stellar characteristics of NP~Aur to those of the other stars in our sample, we find that all stars with TiO emission show large-amplitude pulsations, s-process enrichment, and a low C/O ratio. These characteristics might be necessary requirements for a star to show TiO in emission, but they are not sufficient.} | The spectral class of S stars contains objects that show absorption bands from oxides such as ZrO, LaO and YO in addition to the absorption bands of TiO that characterize M stars \citep{Merrill1922, Keenan1954}. It is often said that S-type stars have C/O ratios close to one. Although the ZrO bands become more pronounced when the C/O ratio is closer to unity, a star can show ZrO absorption bands while the C/O ratio is as low as 0.5 \citep{VanEck2010}. This is only possible if the stars have enhanced abundances of s-process elements such as Zr, La and Y. If this s-process enrichment is caused by nucleosynthesis and dredge-up on the thermally-pulsing AGB, these stars are \emph{intrinsic} S stars. \emph{Extrinsic} S stars are enriched through pollution by a binary companion \citep{Groenewegen1993, VanEck1999}. The infrared spectra of S stars are diverse and show a mix of dust species typical for either oxygen-rich or carbon-rich circumstellar environments, molecular emission bands, or no excess emission at all \citep{Chen1993, Hony2009, Smolders2010, Smolders2012}. This heterogeneity of the infrared appearance within one spectral class is caused by the peculiar chemical composition of the stellar atmosphere of S stars. In oxygen-rich stars, the stable CO molecule forms at high temperatures in the stellar outflows until almost all free carbon is consumed, leaving only the free oxygen atoms to produce molecules and dust. For carbon-rich stars, the oxygen is depleted and only carbon is left to form molecules or dust grains. This CO dichotomy explains the differences in spectral appearance \citep{Millar2000, Willacy1997, Treffers1974}. Because S-type stars span a range of C/O ratios from near solar values to near unity, and since the composition of the molecular shell and dusty wind is highly sensitive to the actual C/O ratio, a rich diversity in gas and solid-state species can be expected \citep{Smolders2012}. The results discussed in this paper are part of a program to study a large sample of S-type AGB stars (Program ID 30737, P.I. S.\,Hony), observed with the Infrared Spectrograph \citep[IRS, ][]{Houck2004} onboard the Spitzer Space Telescope \citep{Werner2004}. The sample was selected from the Stephenson S star catalog \citep{stephenson2}, limited to those targets that are observable with Spitzer based on the 2MASS and IRAS data. The extrinsic S stars were excluded from the sample based on the amplitude of variability \citep{Smolders2012}. The final sample consists of 87 S-type stars in total, containing 32 stars with dust emission features, 3 stars with only molecular emission features, 4 stars with hydrocarbon emission, and 47 stars without significant emission features. In addition to the infrared spectra, we acquired high-resolution optical spectroscopy using the HERMES \'{e}chelle spectrograph at the Mercator telescope in La Palma \citep{Raskin2011}, covering 3770-9000\,\AA\ with a resolving power of $\mathrm{R} \approx 85\,000$. \citet{Smolders2012} provide more details on the data and the sample selection. Previous studies of this sample have already led to the identification of the fundamental ro-vibrational band of SiS at 13\,\um and its first overtone at 6.7\,\um in absorption \citep{Cami2009} and in emission \citep{Sloan2011}. Furthermore, four stars in this sample show the typical hydrocarbon emission features on top of an oxygen-rich photosphere \citep{Smolders2010}. In this letter we focus on NP\,Aurigae, which shows a double peaked emission feature in the 9.8--10.5~\um region, which we identify as the fundamental vibrational band of TiO. In Sects. 2 and 3 we present the identification and modeling of the emission feature. In Sect. 4, we discuss the search for TiO and TiO$_2$ in circumstellar environments and we show the tentative detection of TiO in two other S-type stars. Finally, we conclude in Sect. 5 with the implications of this finding. | We presented the detection of a new, unusual emission feature in the 9-11\,\um region. We identified this feature as the fundamental vibrational band of gaseous TiO. We showed that it is possible to reproduce this emission band using a simple isothermal, single molecular slab. This model can constrain the temperature of the TiO molecules, but because the emission comes from an approximately optically thin region, we cannot put strict constraints on the column density or outer radius of the emitting region. We discussed that this is the first time that TiO molecules are observed in the circumstellar environment of an AGB star. Based on the stellar properties of NP~Aur, the two stars with a tentative detection of this TiO emission band, and the other S stars in our sample, we argue that (i) large-amplitude pulsations, (ii) a low C/O ratio, and (iii) weak dust emission in the 9-11\,\um region might be necessary, but not sufficient requirements for the infrared spectrum to show TiO emission. | 12 | 6 | 1206.2451 |
1206 | 1206.5792_arXiv.txt | In this study, we discuss stellar spots, stellar flares and also the relation between these two magnetic proccess that take place on UV Ceti stars. In addition, the hypothesis about slow flares described by \citet{Gur86} will be discussed. All these discussions are based on the results of three years of observations of the UV Ceti type stars AD Leo, EV Lac, V1005 Ori, EQ Peg and V1054 Oph. First of all, the results show that the stellar spot activity occurs on the stellar surface of EV Lac, V1005 Ori and EQ Peg, while AD Leo does not show any short-term variability and V1054 Oph does not exhibits any variability. We report new ephemerides, for EV Lac, V1005 Ori and EQ Peg, obtained from the time series analyses. The phases, computed in intervals of 0.10 phase length, where the mean flare occurence rates get maximum amplitude, and the phases of rotational modulation were compared to investigate whether there is any longitudinal relation between stellar flares and spots. Although, the results show that flare events are related with spotted areas on the stellar surfaces in some of the observing seasons, we did not find any clear correlation among them. Finally, it is tested whether slow flares are the fast flares occurring on the opposite side of the stars according to the direction of the observers as mentioned in the hypothesis developed by \citet{Gur86}. The flare occurence rates reveal that both slow and fast flares can occur in any rotational phases. The flare occurence rates of both fast and slow flares are varying in the same way along the longitudes for all program stars. These results are not expected based on the case mentioned in the hypothesis. | Many samples of UV Ceti type stars posses the stellar spot activity known as BY Dra Syndrome. The BY Dra Syndrome among the UV Ceti stars was first found by \citet{Kro52}. He reported the existence of sinusoidal-like variations at out-of-eclipses of the eclipsing binary star YY Gem. \citet{Kro52} explained this sinusoidal-like variation at out-of-eclipse as a heterogeneous temperature on star surface, which was called BY Dra Syndrome by \citet{Kun75}. This interpretation of BY Dra Syndrome in terms of dark regions of the surface of rotating stars was confirmed, based on more rigorous arguments, by later works of \citet{Fer73, Bop73, Vog75, Fri75}. Since the most of solar flares occur over solar spot regions, in the stellar case it is also expected to find a correlation between the frequency of flares and the effects caused by spots in the light curve. In order to determine a similar relation among the stars, lots of studies have been made using UV Ceti type stars showing stellar spot activity such as BY Dra. One of these studies was reported by \citet{Bop74} on YY Gem. \citet{Bop74} did not find a clear correlation between the location and extent in longitude of flares and spots on YY Gem. Moreover, he notes that the longitudinal extent he derived for a flare-producing region is in good agreement with the longitudinal extents of starspots previously calculated for BY Dra and CC Eri by \citet{Bop73}. In another study, \citet{Pet83} compared longitudes of the stellar spots obtained from two years observations of YZ CMi and EV Lac with longitudes of flare events and distributions of flare energy and frequencies along the longitude obtained in the same work. Direct comparisons and statistical tests are not able to reveal positive relationships between flare frequency or flare energy and the position of the spotted region. In another extended work, \citet{Let97} looked for whether there is any relation between stellar spots and flares observed from 1967 to 1977 in the observations of EV Lac. The authors were able to find a relation in the year 1970. They could not find any relation in other observing seasons because of the higher threshold of the system used for flare detection. In the last years, \citet{Gar03} found some flares occurring in the same active area with other activity patterns with using simultaneous observations. Since no correlation is found between stellar spots and flares, a hypothesis about fast and slow flares was put forward. The hypothesis is based on the work named as Fast Electron Hypothesis. According to this hypothesis, the shape of a flare light variation depends on the location of the event on the star surface in respect to direction of observer. If the flaring area is on the front side of the star according to the observer, the light variation shape looks as a fast flare. If the flaring area is on the opposite side of the star according to the observer, the light variation shape looks as a slow flare \citep{Gur65, Gur86}. In addition, \citet{Gur88} described two types of flares to model flare light curves. \citet{Gur88} indicated that thermal processes are dominant in the processes of slow flares, which are 95$\%$ of all flares observed in UV Ceti type stars. Non-thermal processes are dominant in the processes of fast flares, which are classified as "other" flares. According to \citet{Gur88}, there is a large energy difference between these two types of flares. Moreover, \citet{Dal10} developed a rule to the classifying of fast and slow flares. When the ratios of flare decay times to flare rise times are computed for two types of flares, the ratios never exceed 3.5 for all slow flares. On the other hand, the ratios are always above 3.5 for fast flares. It means that if the decay time of a flare is 3.5 times longer than its rise time at least, the flare is a fast flare. If not, the flare is a slow flare. In this paper, the results obtained from Johnson UBVR observations of AD Leo, EV Lac and V1005 Ori will be discussed. \citet{Sha74} and \citet{Bop77} reported that V1005 Ori is a flare star that exhibits rotational modulation due to stellar spots. The authors found an amplitude variation of $0^{m}.08$ with a period of $1^{d}.96$ in V band. Besides, \citet{Bop78} examined photometric data in the time series analyses and found 4 periods for rotational modulation. The period of $1^{d}.858$ is suggested as the most probable period among them. On contrary, in B band observations of 1981, a $4.56\pm0.01$ day period variation with an amplitude of $0^{m}.16$ was found \citep{Byr84}. No important light curve changes are seen in the years 1996 and 1997, while the minimum phases of rotation modulation is varied from $0^{P}.40$ to $0^{P}.55$. The amplitude of the curves is $0^{m}.10$ in the year 1996, but in the year 1997 it gets larger than the previous ones \citep{Ama01}. In the case of AD Leo, it is a debate issue whether AD Leo has any stellar spot activity, or not. \citet{Chu74} and \citet{Mul74} show that AD Leo does not exhibit any rotational modulation caused by stellar spots. Besides, \citet{And79} found no variations at the $0.02$ magnitude level during the period 1978, May 10 to 17. However, \citet{Spi86} reveals that AD Leo demonstrates BY Dra Syndrome with a period of $2.7\pm0.05$ days. In addition, \citet{Pan93} confirmed this period of AD Leo for BY Dra variation. On the other hand, EV Lac is a well known active star with both high level flare and stellar spot activities. \citet{Mah81} indicated that the star has no variation caused by rotational modulation in the observations in B band from 1972 to 1976. However, \citet{Pet80} showed a rotational modulation with a period of $4^{d}.378$ and an amplitude of $0^{m}.07$. \citet{Pet83}, based on continuous observations from 1979 to 1981, renewed the ephemerides of the variation as a period of $4^{d}.375$ and an amplitude of $0^{m}.08$. This indicates that the light curves of EV Lac were almost constant for 2.5 years because the spot groups on the star are stable during these 2.5 years. Using the renewed ephemerides, \citet{Kle87} found the amplitude of light curve enlarging from $0^{m}.08$ to $0^{m}.16$ in the year 1986. On the other hand, no variation was seen in the light curve of the year 1987. \citet{Pet92} showed that the spotted area is located in the same semi-sphere on EV Lac for 10 years. Comparing the phases of the light curve minima caused by rotational modulation with the flare frequencies and the distribution of the flare equivalent durations for YZ CMi and EV Lac, \citet{Pet83} showed that there is no relation between the flare activity and stellar spot activity on these stars. EQ Peg is classified as a metal-rich star and it is a member of the young disk population in the galaxy \citep{Vee74, Fle95}. EQ Peg is a visual binary \citep{Wil54}. Both components are flare stars \citep{Pet83}. Angular distance between components is given as a value between 3$\arcsec$.5 and 5$\arcsec$.2 \citep{Hai87, Rob04}. One of the components is 10.4 mag and the other is 12.6 mag in V band \citep{Kuk69}. Observations show that flares on EQ Peg generally come from the fainter component \citep{Fos95}. \citet{Rod78} proved that 65$\%$ of the flares come from faint component and about 35$\%$ from the brighter component. The fourth star in this study is V1054 Oph, whose flare activity was discovered by \citet{Egg65}. \citet{Nor07} demonstrated that EQ Peg has a variability with the period of $1^{d}.0664$. V1054 Oph (= Wolf 630ABab, Gliese 644ABab) is a member of Wolf star group \citep{Joy47, Joy74}. Wolf 630ABab, Wolf 629AB (= Gliese 643AB) and VB8 (= Gliese 644C), are the members of the main triplet system, whose scheme is shown in Fig.1 given by \citet{Pet84}. The masses were derived for each components of Wolf 630ABab by \citet{Maz01}. The author showed that the masses are 0.41 $M_{\odot}$ for Wolf 629A, 0.336 $M_{\odot}$ for Wolf 630Ba and 0.304 $M_{\odot}$ for Wolf 630Bb. In addition, \citet{Maz01} demonstrated that the age of the system is about 5 Gyr. In this study, for each program stars, we analyse the variations at out-of-flare for each light curves obtained in Johnson UBVR observations, or not. Although all of them show high flare activity, EV Lac, V1005 Ori and EQ Peg exhibit stellar spot activity. On the other hand, the spot activity is not obvious for AD Leo. It is discussed whether AD Leo has any stellar spot activity, or not. Finally, this work do not demonstrate any variation from rotational modulations. To perform this kind of studies we would require a long term observing program. As a part of this study, the phase distributions for both fast and slow flares are examined in terms of the minimum phases of rotational modulation. Thus, hypothesis developed by \citet{Gur86} is tested. | \subsection{Stellar Rotational Modulation and Stellar Spot Activity} Most of the UV Ceti type stars are full convective red dwarfs with sudden-high energy emitting. As it can be seen in the literature, BY Dra Syndrome at out-of-flares is seen in a few stars among 463 flare stars catalogued by \citet{Ger99}. EV Lac and V1005 Ori can be given as two examples because the studies in the literature and this study indicate that both stars show the variation due to rotational modulation at out-of-flares. In the case of EV Lac, the time series analyses show that the period of rotational modulation found for each data set is range from $4^{d}.330$ to $4^{d}.378$. The periods found are similar to those found by \citet{Pet83} and \citet{Mah81}. Although the periods found for each season are a little bit different, this difference is relatively small. When the amplitudes of the light curves are examined for EV Lac, the amplitude of this variation was dramatically decreasing from the year 2004 to 2005, while the amplitude was clearly larger than ever in this study. However, the mean average of brightness in the light curves was slowly decreasing from the year 2004 to 2006. The minima phases of the light curves for the three seasons were computed and, it was found as $0^{P}.62$ for the season 2004, $0^{P}.54$ for 2005 and $0^{P}.60$ for 2006. In the case of V1005 Ori, the periods of the rotational modulation for each season are range from $4^{d}.419$ to $4^{d}.429$. In the literature, \citet{Bop78} found four possible periods varied from $1^{d}.883$ to $2^{d}.199$. On the other hand, \citet{Byr84} found a period of $4^{d}.565$. As it is seen, the periods found in this study are close to the period found by \citet{Byr84}. When the amplitudes of the light curves were examined, the amplitude observed in the season of 2004/2005 was so smaller than the ones observed in the previous and later seasons that there was no minimum in the light curve. Although the mean average of brightness in the light curves was not changing, the minimum phases of the light curves were varying. The minimum phase of the curve for the season 2004/2005 was about $0^{P}.34$ and about $0^{P}.66$ for the season 2006/2007. It is hard to say that the minimum phase of the light curves for the season 2005/2006 was about $0^{P}.78$. The case of AD Leo is different from the other two stars. The time series analyses do not show any regular variation over the $3\sigma$ level in one season. On the other hand, the mean brightness levels were increased a value of $0^{m}.01$ from the first season to the second and a value of $0^{m}.02$ from the second to the last season. This can be because of the stellar polar spots. If the literature is considered, the stellar spots can be carried to polar regions in the case of rapid rotation in the young stars \citep{Sch92, Sch96}. According to \citet{Mon01}, AD Leo is at the age of 200 Myr. The range of equatorial rotational velocity ($vsini$) given in the literature is between $5$ - $5.8$ and $9.0$ $kms^{-1}$ for AD Leo \citep{Mar92, Pet91}. Besides, considering these values of $vsini$, the real rotational velocities must be larger than these values. If both the age and equatorial rotational velocity value parameters in these papers are considered, according to \citet{Sch92} and \citet{Sch96}, some spots might be located on the polars for AD Leo. In fact, \citet{Pet92} indicate that BY Dra had spotted area near polar region, which was stable for 14 years and EV Lac has a similar area for 10 years. If the studies made by \citet{Spi86} and \citet{Pan93} are considered, AD Leo might sometimes show rotational modulation due to the spotted area occurring near the equatorial regions. On the other hand, there is another probability. If the colour index of V-R is considered, it is seen that the star gets bluer from a season to next one, when the star get brighter. Besides, no amplitude is seen in the light curves. These can be some indicators that all the surface of the star is covered by cool spots and the efficiency of the spots gets weaker from one season to next one. The colour curves of both EV Lac and V1005 Ori sometimes exhibit a clear colour excess around the minimum phases of the light curves for some observing seasons. This can be an indicator of some bright areas such as faculae on the surface of these young stars. The effects of the bright areas such as faculae can be seen in the variations of B-V and sometimes V-R colour, while these effects are not seen in the variations in the light curves of BVR bands due to cool spots. The cool spots are more efficient in the light curves of B, especially V and R bands. The same effect is seen in the variations caused by the flare activity. Although there is some small effects or no effect of the flare activity in V and R bands, but there is some clear variations in U band light and U-B colour curves. In this study, we observed the stars in U band to investigate the variations out-of-flares. In a sense, U band observations is used to control whether there is any flare activity in the observing durations. If there is some variations in U band light and U-B colour, we did not used the observation to investigate the variation out-of-flares. \citet{Nor07} showed that EQ Peg has a variability with the period of $1^{d}.0664$. In this study, the time series analyses supported this period. According to our analyses, EQ Peg exhibits short-term variability with the period of $1^{d}.0608$. However, it is seen that there is not any variability in the colour indexes. Analysing the light curve of EQ Peg, it was found as $0^{P}.32$ for the minimum phase of the rotational modulation. \subsection{The Relation Between Stellar Spot and Flare Activities} There are many studies about whether the flares of UV Ceti type stars showing BY Dra Syndrome are occurring at the same longitudes of stellar spots, or not. Having the same longitudes of flare and spots is an expected case for these stars, because solar flares are mostly occurring in the active regions, where spots are located on the Sun \citep{Ben10}. In the respect of Stellar-Solar Connection, a result of the $Ca$ $II$ $H\&K$ Project of Mount Wilson Observatory \citep{Wil78, Bal95}, if the areas of flares and spots are related on the Sun, the same case might be expected for the stars. In fact, \citet{Mon96} have found some evidence to demonstrate this relations. Besides, \citet{Let97} have found a variations of both the rotational modulation and the phase distribution of flare occurence rates in the same way for the observations in the year 1970. On the other hand, no clear relation between stellar flares and spots has been found by \citet{Bop74, Pet83}. However, \citet{Pet83} did not draw firm conclusions because of being a non-uniqueness problem. In this study, the flare occurence rates, the ratio of flare number to monitoring time, were computed in intervals of 0.10 phase length as the same method used by \citet{Let97} with just one difference. The flare maximum times were used to compute the phases due to main energy emitting in this part of the flare light curves. We observed AD Leo for 79.61 $h$ and detected 119 flares in three seasons. EV Lac was observed 109.63 $h$ and 93 flares were detected in three seasons. V1005 Ori was observed for 44.75 $h$ and 44 flares were detected in two seasons. EQ Peg was observed for 100.26 $h$ and 73 U band flare were detected. Since no rotational modulation was found to compare for AD Leo, all the flares detected in three season were combined in order to just find whether there is any phase, in which the flare occurence rate gets a peak. On the other hand, we examined flare phase distributions for each season for both EV Lac and V1005 Ori. In the case of these stars, if the distribution of flares did not cover almost all phases in an observing season of a star, the season is neglected for the comparison of flare and spot activity. Consequently, for both EV Lac and V1005 Ori, we chose the seasons, in which the best flare distributions were obtained. Thus, we only used the seasons, in which there is enough data to get reliable conclusions about flare occurrence distributions. In addition, to determine the phases of MFOR, all the distributions were modelled with the polynomial function. Resolving these models, maximum flare occurrence rates and their phase were found for all program stars. In the case of EV Lac, no relation is seen between the minimum phase of the rotational modulation and the phase, in which flare activity reaches the MFOR. The minimum phase of the rotational modulation observed in the season 2004 is $0^{P}.62$, while the phase of MFOR is $0^{P}.75$. The minimum phase of rotational modulation is $0^{P}.54$, while the flare occurrence rate reaches maximum level in about the phase of $0^{P}.45$ for the season 2005. In the last season of EV Lac, rotational modulation minimum is seen in $0^{P}.60$, as MFOR is in $0^{P}.35$. In the case of V1005 Ori, there is enough data in only one season to compare. As it is seen, the minimum phase of rotational modulation is $0^{P}.78$, while phase of MFOR is about $0^{P}.87$ for the season 2005/2006. In this study, the time series analyses indicated that AD Leo does not have any rotational modulation. Therefore, any minimum time could not have been determined from the observations of three seasons for AD Leo. Because of this, we could not compare the rotational modulation with flare activity in the case of AD Leo. On the other hand, using combined data of three seasons, we found that the MFOR is seen in $0^{P}.45$. This phases was computed with using the ephemeris given in Equation (1) taken from \citet{Pan93}. The time series analyses do not show any short-term variation in the light curves of AD Leo. Because of this, we waited that there is no any phase, in which the flare activity gets higher levels. On the other hand, as it is seen from the histogram and its Normal Gaussian model for AD Leo, there is a phase for MFOR. Considering the phase of MFOR, the active region(s) in some particular part of the surface can be more active than the others on the surface of the star. Considering the light and colour curves of AD Leo, almost all surface of the star may be covered by stellar spots, while it is seen that some region(s) in the surface of the star can be more active than the remainder of the surface. In the case of EQ Peg, the minimum phase of the rotational modulation is $0^{P}.32$, while the phase of MFOR is $0^{P}.95$. The results acquired from EV Lac and V1005 Ori demonstrated that flare activity can reach high levels at almost the same longitudes, in which stellar spots occur. On the other hand, there is a considerable difference between the phases of stellar spot and MFOR for the observing season 2007 of EV Lac. In conclusion, it is seen that there is a longitudinal relation between stellar spot and flare activities in general manner. Nevertheless, there are some differences and this makes difficult to do a definite conclusion. Moreover, in the case of EQ Peg, the MFOR gets the minimum towards the minimum phase of the rotational modulation. All these cases can be because of a dynamo which is working in the red dwarf stars. In spite of the Sun, red dwarf stars are mostly known to have a different dynamo because of full convective outer atmosphere. However, in the last years, some studies showed that flares on the Sun do not have to be located upon the spotted areas on the Sun \citep{Bor07}. In addition, it should be kept in mind that most of the studies have been done with using the data obtained from white-light flare observations, but a white-light flare does not have to occur in a flare process. Recent studies have shown that non white-light flares may be so common in UV Ceti-type stars as they are in the Sun \citep{Cre04, Cre06}. In this point, it can be mentioned that the analyses of data obtained from only white-light flare observations are not sufficiently qualified. For instance, \citet{Gar03} found some flares occurring in the same active area with other activity patterns with using simultaneous observations. \subsection{Phase Distribution of The Fast And Slow Flares} Using the inverse Compton event, \citet{Gur86} developed a hypothesis called Fast Electron Hypothesis, in which red dwarfs generate only fast flares on their surface. On the other hand, according to the flare region on the surface of the star in respect to direction of observer, the shapes of the flare light variations can be seen like a slow flare \citep{Gur86}. If the scenario in this hypothesis is working, it is expected that the fast and slow flares should collected into two phases in the light curves of UV Ceti type stars showing BY Dra Syndrome. It is also expected that these two phases are separated from each other with intervals of $0^{P}.50$ in phase. In this study, according to the rule described by \cite{Dal10}, the flares are classified as fast and slow flares. Then the phase distributions of fast flares were compared with the phases of slow flares in order to find out whether there is any separation as expected in this respect. When the phases of both fast and slow flares are examined one by one, it is clear that both of them can occur in any phase. To reach a definite result, the phase distributions of both fast and slow flares are statistically investigated. As it is stated in the previous section, if the distribution of flares did not cover almost all phases in an observing season of a star, the season is neglected for that star. Consequently, we chose the seasons, in which there is enough data to get reliable conclusions about flare occurrence distributions for both fast and slow flares. In the case of AD Leo and EQ Peg, we combined all the fast flares of three seasons as we made for the slow flares. For both fast and slow flares, using Equation (5), the number of flares occurring per an hour in intervals of 0.10 phase length was computed. The obtained occurrence rates for both fast and slow flares are shown by histograms in Figures 12, 13, 14 and 15. Once again, all the distributions were modelled with the polynomial function. Resolving these models, maximum flare occurrence rates and their phase of both slow and fast flares were found for all program stars. In the case of AD Leo, the analyses show that both fast and slow flares have a difference of $0^{P}.17$ between the phases, in which flare occurrence rates in intervals of 0.10 phase length reach maximum amplitudes. The same difference is $0^{P}.05$ for EV Lac in the season of 2004. Although these differences are acceptable as low values according to Fast Electron Hypothesis, the difference seen in the season of 2006 is $0^{P}.50$ for EV Lac. This value is the expected value in respect of Fast Electron Hypothesis. In the case of V1005 Ori, slow and fast flares could be compared only for the season of 2005/2006. The result is that both fast and slow flares have a difference of $0^{P}.30$ between the phases of maximum flare occurrence rates. In the case of EQ Peg, the phase difference between MFORs of slow and fast flares is about $0^{P}.40$. The value obtained from EQ Peg is also the expected value in respect of Fast Electron Hypothesis. It should be noted that in the case of EQ Peg, it is seen just one clear peak for the distribution of MFOR for the fast flares, while there are several peaks for the slow flares. As it is seen from the analyses, both the fast and the slow flares sometimes the same longitudinal distributions and sometimes different. This makes difficult to say that there is a regular longitudinal division between these two types of flares as expected according to \citet{Gur86}. This means that, when a slow flare is observed, it does not have to be a fast flare occurred on the opposite side of the star in respect to observer direction. | 12 | 6 | 1206.5792 |
1206 | 1206.4382_arXiv.txt | Despite expanding research activity in gravitational lens modeling, there is no particular software which is considered a standard. Much of the gravitational lens modeling software is written by individual investigators for their own use. Some gravitational lens modeling software is freely available for download but is widely variable with regard to ease of use and quality of documentation. This review of 13 software packages was undertaken to provide a single source of information. Gravitational lens models are classified as parametric models or non-parametric models, and can be further divided into research and educational software. Software used in research includes the GRAVLENS package (with both gravlens and lensmodel), Lenstool, LensPerfect, glafic, PixeLens, SimpLens, Lensview, and GRALE. In this review, GravLensHD, G-Lens, Gravitational Lensing, lens and MOWGLI are categorized as educational programs that are useful for demonstrating various aspects of lensing. Each of the 13 software packages is reviewed with regard to software features (installation, documentation, files provided, etc.) and lensing features (type of model, input data, output data, etc.) as well as a brief review of studies where they have been used. Recent studies have demonstrated the utility of strong gravitational lensing data for mass mapping, and suggest increased use of these techniques in the future. Coupled with the advent of greatly improved imaging, new approaches to modeling of strong gravitational lens systems are needed. This is the first systematic review of strong gravitational lens modeling software, providing investigators with a starting point for future software development to further advance gravitational lens modeling research. | Gravitational lensing has great promise to provide new insights into the structure and history of the universe. Gravitational lensing has yielded many exciting results by mapping dark matter distributions, and the recent use of strong gravitational lensing data has added a new dimension to this research \citep{Coe2010}. Gravitational lensing is a very active area of investigation, and research is highly dependent on computer modeling. In some areas of contemporary astrophysics research, there is software that is a \emph{de facto} standard for many investigators (e.g. SExtractor \citep{Sextractor} \footnote{\url{http://astroa.physics.metu.edu.tr/MANUALS/sextractor/}}, GALFIT \citep{Galfit} \footnote{\url{http://users.obs.carnegiescience.edu/peng/work/galfit/galfit.html}}, etc.). While gravitational lens modeling software has been written, there are no standards and no easily accessible source of information about existing software. The lack of standard software may be a virtue of the gravitational lensing community, allowing more flexibility and creativity. The lack of a single standard program makes it more important to compare existing software used for modeling strong gravitational lenses. Information regarding existing software will be helpful to those developing new approaches and interfaces. This review was undertaken to identify available strong lens modeling software, and review the installation, use, and the nature of inputs and data outputs. This paper serves as a guide to available software and provides useful information to both new and established investigators in this field. The availability of source code may be a useful starting point for anyone writing their own modeling software. This paper is organized as follows. In section \S \ref{Classif}, we review the classification of gravitational lens models and the methodology used to review the available software. In section \S \ref{Res}, we review eight software packages that have been used extensively in gravitational lens research. Following this, in section \S \ref{Educ}, we review five programs that are useful in education for gravitational lensing. In section \S \ref{Compare} we discuss several factors of importance in selecting and comparing available software and in section \S \ref{Concl} we make suggestions for the next generation of software to support future gravitational lens research based on this review. | \label{Concl} There is a wide variety of gravitational lensing modeling software available. Many of the publications using these packages are written by the software developers, suggesting that the software is developed for their personal use. Software available only as executable files has the advantage of being rapidly usable, as long as the computing platform is available. Software distributed as source code may require significant time to compile and prepare given the vagaries of software libraries. The comprehensiveness of this study is limited by the ability to identify existing software or studies related to the software reviewed. However, a fairly wide spectrum of software was available for review, and this analysis identifies opportunities for improvement based on existing software. Awareness of available software may limit the need to develop proprietary software in the future. The source code is available for some applications which facilitates the development of custom software. Given the increased activity in gravitational lensing research, sharing of software and algorithms may result in significant time savings. The usability of available software is somewhat limited by the essentially consistent use of character based user interfaces. Future software should be modular in nature and use a graphical user interface for improved functionality. Until recently, the construction of gravitational mass maps to detail dark matter distribution has depended on data from weak lensing. However, as shown in recent studies using LensPerfect, glafic and GRALE, accurate mass maps can be constructed using data from strong gravitational lensing data, and this may signal a new era in studies of strong gravitational lensing. New approaches to software development will be necessary to support this shift in research, especially with the advent of far more detailed images from the next generation of telescopes. Future studies should include direct comparisons with other available software in addition to indirect comparisons with previous studies of well-described lensing systems. | 12 | 6 | 1206.4382 |
1206 | 1206.3331_arXiv.txt | {The Core Mass Functions of low-mass star-forming regions are found to resemble the shape of the Initial Mass Function (IMF). A similar result is observed for the dust clumps in high-mass star forming regions, although at spatial scales of clusters that do not resolve the substructure that is found in these massive clumps. The region \mia\ is one exception, having been observed at spatial scales on the order of $\sim 2500$\,AU, sufficient to resolve the clump substructure into individual cores.} {We investigate the protostellar content of \mia\ at high spatial resolution at \rxun, determining the temperature structure of the region and deriving its Core Mass Function.} {The massive star-forming region \mia\ was mapped with the PdBI (BCD configurations) at \rxun\ and \rxtres\ in the continuum and several transitions of formaldehyde (\hdco) and methyl cyanide (\cian). The \hdco\ transitions were also observed with the IRAM 30\,m Telescope.} {We detected 26 continuum sources at \rxun\ with a spatial resolution down to $\sim 2200$\,AU, several of them with counterparts at NIR and MIR wavelengths, distributed in two (proto)clusters. With the PdBI \cian\ and PdBI/IRAM 30\,m \hdco\ emission we derived the temperature structure of the region, ranging from 35 to 90\,K. Using these temperatures we calculated the core masses of the detected sources, ranging from $\sim 0.7$ to $\sim 8\,$M\sun. These masses were strongly affected by the spatial filtering of the interferometer, filtering out a common envelope with $\sim 90\%$ of the single-dish flux. Considering only the detected dense cores, and accounting for binning effects as well as cumulative distributions, we derived a Core Mass Function, with a power-law index $\beta=-2.3\pm 0.2$. We resolve the Jeans length of the (proto)clusters by one order of magnitude, and only find little velocity dispersion between the different subsources.} { Since we cannot unambiguously differentiate protostellar and prestellar cores, the derived CMF is not prestellar. Furthermore, because of the large fraction of missing flux, we cannot establish a firm link between the CMF and the IMF. This implies that future high-mass CMF studies will require to complement the interferometer continuum data with the short spacing information, a task suitable for ALMA. We note that the method of extracting temperatures using \hdco\ lines becomes less applicable when reaching the dense core scales of the interferometric observations because most of the \hdco\ appears to originate in the envelope structure. } | \label{sec-intro} Our understanding of the structure of the cold, dense interstellar medium (ISM) in star-forming regions has improved in the last years. In those regions the ISM exhibits a clumpy, often filamentary structure with density maxima at the sites of star formation. To represent this structure quantitatively we use the Core Mass Function (CMF). In this paper we will refer to ``core'' as the small (diameter ${\rm D}\sim 0.01$\,pc), dense condensation that will form individual stars or small multiple systems, while with ``clump'' we denote structures that may form (proto)clusters and may therefore be more massive and larger than cores. In this sense, cores may be considered as a subset of clumps. Sub-mm observations of low-mass star-forming regions such as Serpens (e.g., \citealt{testi1998}), Orion B (e.g., \citealt{motte2001}), Aquila (e.g., \citealt{konyves2010}) and $\rho$ Oph (e.g., \citealt{motte1998}), as well as near-infrared extinction maps (e.g., \citealt{alves2007}), show that their CMFs resemble the shape and intrinsic mass scale of the stellar initial mass function (IMF; e.g., \citealt{salpeter1955,kroupa2002}). This suggests that these dense cores would be the immediate precursors of stars, and that by applying a more or less constant core-to-star mass conversion efficiency we can obtain the IMF from the CMF. In the case of Massive Star-Forming (MSF) regions we have, for example, the analysis of \citet{reid2006}. They gathered the published masses of the MSF regions M8 (e.g, \citealt{tothill2002}), M17 (e.g., \citealt{reid2006}), NGC\,7538 (e.g., \citealt{reid2005}), W43 (e.g., \citealt{motte2003}) and RCW\,106 (e.g., \citealt{mookerjea2004}), tracing spatial scales of clumps that correspond to (proto)clusters rather than to individual cores, and tested the fit of several functional forms for their clump mass functions. They found that in those cases, the best fit was obtained by a double power law, having a mean power-law exponent for the high-mass end consistent with the Salpeter IMF. This would again imply that by an almost one-to-one mass conversion efficiency the IMF could be obtained from the clump mass function, as in the case for low-mass star-forming regions. Similar analysis were conducted by e.g., \citet{shirley2003,beltran2006}. Furthermore, a theoretical study by \citet{chabrier2010} on the relationship between the CMF and the IMF suggests a tight correlation between the two. Once more, this implies that the IMF would be defined by the CMF in the early stages of evolution. However, this one-to-one relationship may not hold as, for example, some clumps must fragment to produce the observed quantity of multiple stellar systems \citep{goodwin2007}. The relatively large distances ($\gtrsim 2\,$kpc) of most of the known MSF regions require a spatial resolution of about $1''$ to resolve the clumps into cores with sizes below $\sim 0.1$\,pc. That resolution in the (sub)mm regime is only achievable with the interferometric technique. So far, only a few MSF regions have been observed in the (sub)mm with spatial resolutions good enough to resolve individual cores (e.g., \citealt{bontemps2010}, \citealt{fontani2009}, \citealt{rodon2008}, \citealt{rathborne2008}, \citealt{beuther2006c}), and only for one source, \mia, has it been possible to determine a core mass function \citep{beuther2004c}. The young MSF region \mia\ is at a distance of $2.16$\,kpc \citep{xu2009} and has an integrated bolometric luminosity of about $10^4$\,L\sun. It is a very active star-forming site, with sources detected from X-rays down to radio wavelengths. It has ${\rm H_2O\ and\ CH_3OH}$ maser emission \citep{sridharan2002,beuther2002c} and X-ray sources \citep{beuther2002e}, denoting the ongoing formation of intermediate-to-high mass stars. The region is embedded within a cluster of over 800 components detected at NIR wavelengths \citep{martin2008,qiu2008}, and it has a rich and energetic outflow component with multiple outflows detected in CO \citep{beuther2002b,beuther2003}. \citet{beuther2002a} found that the large-scale mm emission shows two massive gas clumps roughly aligned in a north-south direction that splits into several subsources with increasing spatial resolution (\citealt{beuther2004c}). With their studies of the mm continuum at high spatial resolution, \beuthert\ were able to derive the mass function of \mia, resulting in a Salpeter-like distribution. However, the strongest caveat in the derivation of that mass function was the fact that a uniform dust temperature was used in the calculation of the masses. Although they argue that the dust temperature distribution should not vary strongly, they also warn that changes in the temperature of the cores would result in a somewhat flattened slope. We have revisited \mia\ observing the mm continuum at high-spatial resolution and obtaining molecular-line emission of known temperature tracers, to determine via several methods a temperature structure for it and in the end derive a more robust mass function. | We resolve the two clumps of the MSF region \mia\ into 26 cores at \rxun, with a spatial resolution of $\sim 2200$\,AU. This resolves the Jeans length of the clumps, and the relative distances between the cores are similar or smaller than the Jeans length corresponding to individual cores. Also, the cores show only a marginal signature of velocity dispersion, implying that the (proto)clusters are not in a strong dynamical evolution. The approach of calculating temperatures using \hdco\ lines might not be suitable for studies at high-spatial resolution and density like this one. We find that the \hdco\ and the continuum emission only correlate for the brightest source in each (proto)cluster and not for the fainter sources, therefore preventing the derivation of the temperature of each single core. However, it allows to derive a temperature structure for the (proto)clusters as a whole. The temperature structure of \mia\ was determined from its \hdco\ and \cian\ emission allowing the estimation of a CMF. Taking into account the arbitrariness of the mass binning when deriving a mass function, we found a CMF index $\beta=-2.3\pm0.2$ for core masses above $\sim 1.4$\,M\sun, confirming the previous results of \beuthert\ with increased confidence levels. It was not possible to determine with enough confidence the evolutionary stage of all the sources detected (i.e, whether they are pre-, proto-, or stellar sources). Because of this, the CMF we derive is not a fully pre-stellar one, and therefore it cannot be considered as a true predecessor of the stellar IMF, but a step in between a pre-stellar CMF and the IMF. An additional important caveat is the difficulty to combine continuum single-dish and interferometer data because of the vastly different bandpasses. This induces missing flux problems severely affecting the mass estimates. To the present day, mapping of the CMF has been done mostly in low-mass star forming regions. We have shown and discussed the caveats that involve such a study in MSF regions, which in turn explains the rarity of such studies so far. Mapping a MSF region down to $ \sim0.01 $\,pc scales requires an angular resolution on the order of one arc-second for the closest regions, and tenths of arc-second if we go to larger distances. At (sub)mm wavelengths that can only be achieved with interferometry. The interferometers operating at this moment normally achieve $ \sim1'' $ resolution, but to go below it is not that common. Sensitivity can become also an issue, and interferometers naturally filter out a large part of the incoming flux. When fully operational, ALMA will improve significantly on these observational limitations. Despite the difficulties and assumptions that have to be made due to these limitations, the study of the Core Mass Function of high-mass star forming regions has great usefulness for the study of the origin of the IMF. It is already seen that for low-mass stars, the IMF and the CMF are indistinguishable, suggesting the kind of relationship between them. That is yet unknown for high-mass stars and therefore it needs to be addressed, both theoretically and observationally. | 12 | 6 | 1206.3331 |
1206 | 1206.2101_arXiv.txt | We review Gamma-Ray Burst (GRB) afterglow follow-up observations being carried out by our group in Korea. We have been performing GRB follow-up observations using the 4-m UKIRT in Hawaii, the 2.1-m telescope at the McDonald observatory in Texas, the 1.5-m telescope at Maidanak observatory in Uzbekistan, the 1.8-m telescope Mt. Bohyun Optical Astronomy Observatory (BOAO) in Korea, and the 1.0-m remotely operated telescope in Mt. Lemmon, Arizona. We outline our facilities, and present highlights of our work, including the studies of high redshift GRBs at $z > 5$, and several other interesting bursts. | GRBs are the most energetic explosions in the universe which result from deaths of stars in special ways. Their enormous energy output ($E_{iso} \sim 10^{53} \, erg\, s^{-1}$) produces very bright afterglows in UV/optical/NIR which can be followed up using moderate-sized ground-based telescopes to understand the nature and the origin of GRB and the emission mechanisms of the afterglow. The bright afterglows serve as a useful tool to understand the distant universe since they could be observed beyond $z > 10$, deep into the re-ionization epoch. They offer valuable information on the re-ionization state and the dust origin in the early universe, as well as the metal enrichment history over the cosmic time. In Korea, we started a GRB follow-up observation from 2007 October \cite{Lee10}. Initially, we focused on the study of the early universe using GRBs as bright probes, but the program evolved over the time, expanding the facilities and the scientific scope of our studies. Here, we describe our GRB follow-up observation activities. | 12 | 6 | 1206.2101 |
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1206 | 1206.0618_arXiv.txt | \noindent Spherical collapse has turned out to be a successful semi-analytic model to study structure formation in different DE models and theories of gravity, but nevertheless, the process of virialization is commonly studied on the basis of the virial theorem of classical mechanics. In the present paper, a fully generally-relativistic virial theorem based on the Tolman-Oppenheimer-Volkoff (TOV) solution for homogeneous, perfect-fluid spheres is constructed for the Einstein-de Sitter and $\Lambda$CDM cosmologies. We investigate the accuracy of classical virialization studies on cosmological scales and consider virialization from a more fundamental point of view. Throughout, we remain within general relativity and the class of FLRW models. The virialization equation is set up and solved numerically for the virial radius, $y_\mathrm{vir}$, from which the virial overdensity $\Delta_V$ is directly obtained. Leading order corrections in the post-Newtonian framework are derived and quantified. In addition, problems in the application of this formalism to dynamical DE models are pointed out and discussed explicitly. We show that, in the weak field limit, the relative contribution of the leading order terms of the post-Newtonian expansion are of the order of $10^{-3}\%$ and the solution of Wang \& Steinhardt (1998) (see \cite{wang_cluster_1998}) is precisely reproduced. Apart from the small corrections, the method could provide insight into the process of virialization from a more fundamental point of view. | \label{int} The question of how structures form in the Universe is a long-standing topic in theoretical cosmology and provides a lot of room for discussion. Since the fully non-linear regime cannot be accessed analytically, huge N-body simulations have been set up to describe structure formation by gravitationally interacting particles in an expanding background. However, these attempts are computationally costly and therefore perturbative approaches have been developed in order to keep the continuous character of general relativity and the FLRW model and make use of methods from fluid mechanics. A very simple semi-analytic model of this kind is the \textbf{Spherical Collapse}. A spherical, overdense patch evolves with the background expanding universe, slows down due to its self-gravity, turns-around and collapses. The object is stabilized by virialization which prevents it from collapsing into a singularity. Despite its simplicity and idealizations, this model gives a first insight into the formation of spherical halos at all mass scales. The underlying formalism dates back to Gunn and Gott in 1972 (see \cite{gunn_infall_1972}), but has been rediscovered and continuously extended in recent years (see \cite{abramo_structure_2007, bartelmann_non-linear_2006, basilakos_virialization_2007, lee_spherical_2010, manera_cluster_2006, maor_spherical_2007, maor_virialization_2005, mota_spherical_2004, nunes_structure_2006, pace_spherical_2010, wang_cluster_1998, wang_virialization_2006, wintergerst_clarifying_2010}).\\ In this work, we are going to use the results of Pace et al. (2010) (see \cite{pace_spherical_2010}) to investigate the process of virialization and try to find answers to some remaining questions in this context. The virial theorem provides a powerful tool to study systems in equilibrium, but in order to clarify its role in the framework of general relativity, a relativistic version is needed. After giving an overview of the classical concepts and the requirements of relativistic calculations in Sects. \ref{sec:class} and \ref{sec:Req}, we will derive a relativistic version of the virial theorem based on the Tolman-Oppenheimer-Volkhoff equation (see \cite{oppenheimer_massive_1939, tolman_static_1939} for the original references) in an Einstein-de Sitter and $\Lambda$CDM universe (see Sects. \ref{sec:TOV_L}, \ref{sec:vir_rel}). In the following, this will be applied to the virialization equation in the spherical collapse model and a post-Newtonian expansion will be performed (see Sects. \ref{sec:vir_eq}, \ref{sec:PN}). The relativistically corrected results for the virial radius and virial overdensity will be discussed and leading order corrections are worked out in particular. (Sect. \ref{sec:res}). We will also dedicate a section to the problems occurring when this formalism is applied to general DE models and point out possible ideas to solve them (see Sect. \ref{sec:dyn_DE}). Throughout the paper, we will make use of natural units, i. e. $c=1$. | We propose a way to set up a fully relativistic method to obtain the virial radius and the virial overdensity for the EdS and $\Lambda$CDM cosmology. Within the assumption of an approximately time-like Killing vector field of the FLRW metric, static solutions for perfect fluid spheres in general relativity (namely the Tolman-Oppenheimer -Volkoff equation) have been successfully applied to extend the virial theorem in a consistent manner. The result has been inserted into the virialization equation which can be solved for the virial radius to find the corresponding virial overdensity. It turns out that the solution of Wang \& Steinhardt (1998) (\cite{wang_cluster_1998}) for the virial radius in a $\Lambda$CDM cosmology is almost perfectly reproduced by our formalism which can also be shown analytically by performing the weak field limit. The first order post-Newtonian expansion has been investigated and the leading order corrections have been worked out and calculated numerically. We found out that they have a relative contribution of $10^{-3}\%$ with respect to the classical term which is of very small, but expected size. Although these corrections are of limited astrophysical interest, the concept itself is a small step towards a more fundamental understanding of virialization of spherical halos in the presence of DE. In addition, an iterative method has been set up to calculate the exact virial redshift numerically. The results of Lee \& Ng (2010) (\cite{lee_spherical_2010}) are reproduced extremely well.\\ Naturally, galaxies and galaxy clusters are far more complex than homogeneous and isotropic spheres, but the spherical collapse model provides a very simple semi-analytic method that already suffices to estimate important parameters like the virial radius and virial overdensity. The process of virialization itself is an additional condition that has been introduced to prevent a spherical overdensity of pressureless dark matter from collapsing into a singularity. A pressure profile, which a relaxed continuous spherical object must obey due to general relativity, can provide new insight into the process of virialization itself.\\ There are certainly some topics this paper cannot address, because they are far beyond its scope. Although the spherical collapse models are powerful tools to obtain estimates of the evolution of structures in the universe, they are limited by their simplicity. In particular, the fact that the spherical overdensity is in no way embedded continuously into the background Friedmann universe is still very idealized and dissatisfying. Secondly, DE cannot be described yet in a self-consistent way with general relativity, since local energy-momentum-conservation, which is required by the theory, is not fulfilled and a coordinate representation of the two fluids is missing. Approaches based on the Lema\^{i}tre-Tolman-Bondi models (see \cite{bondi_spherically_1999, marra_exact_2012, pereira_evolution_2010, romano_constructing_2011, tolman_effect_1997, valkenburg_exact_2011}) as well as the presented work of Misner and Sharp (1964) (see \cite{misner_relativistic_1964}) are promising candidates for following investigations. | 12 | 6 | 1206.0618 |
1206 | 1206.0873_arXiv.txt | The mass distribution of prestellar cores is obtained for clouds with arbitrary internal mass distributions using a selection criterion based on the thermal and turbulent Jeans mass and applied hierarchically from small to large scales. We have checked this methodology comparing our results for a lognormal density PDF with the theoretical CMF derived by Hennebelle \& Chabrier, namely a power-law at large scales and a log-normal cutoff at low scales, but our method can be applied to any mass distributions representing a star-forming cloud. This methodology enables us to connect the parental cloud structure with the mass distribution of the cores and their spatial distribution, providing an efficient tool for investigating the physical properties of the molecular clouds that give rise to the prestellar core distributions observed. Simulated fBm clouds with the Hurst exponent close to the value $H=1/3$ give the best agreement with the theoretical CMF derived by Hennebelle \& Chabrier and Chabrier's system IMF. Likewise, the spatial distribution of the cores derived from our methodology show a surface density of companions compatible with those observed in Trapezium and Ophiucus star-forming regions. This method also allows us to analyze the properties of the mass distribution of cores for different realizations. We found that the variations in the number of cores formed in different realizations of fBm clouds (with the same Hurst exponent) are much larger than the expected root ${\cal N}$ statistical fluctuations, increasing with $H$. | Is the stellar Initial Mass Function (IMF) universal? This question has been in the literature for a long time, and is now extended to the core mass function (CMF) since a close relation between the IMF and the CMF has been recognized (Motte et al. 1998; Testi \& Sargent 1998; Alves et al. 2007; Chabrier \& Hennebelle 2010; Michel et al. 2011). Compared to the CMF, the mass function of stellar systems seems to be shifted to lower masses by a factor that does not depend on the core mass. The currently favored conversion efficiency value of the progenitor core mass to the stellar system is $\sim 1/3$. However, the origin of this conversion efficiency is still controversial \citep{Ada96,Mat00,Eno08,Dib11}. Significant variations in the mass function of young clusters are observed in the disk of the Galaxy \citep[e.g.,][]{Sca98}, but most of these variations are consistent with random sampling from a universal IMF \citep{Elm97,Elm99,Kro02,Bas10,Par11}. The non-lineal processes involved in the star formation process determine on the one hand the universal form of the IMF and on the other hand the range of expected variations of the mass function around this universal form. These fluctuations arise naturally in IMF models based on deterministic chaos \citep{San99} and are also observed in three-dimensional hydrodynamic simulations. Recently \citet{Gir11} performed a parameter study of the fragmentation properties of collapsing isothermal gas cores with different initial conditions and showed that the density profile strongly determines the number of formed stars, the onset of star formation, the stellar mass distribution, and the spatial stellar distribution. Furthermore, the random setup of the turbulent velocity field in SPH simulation has a major impact in the different morphology of the filamentary structure, and consequently on the number of sink particles \citep[as shown by][]{Gir11,Gir12a,Gir12b}. The ever increasing resolution of magnetohydrodynamic numerical simulations will provide the answer to many of these questions \citep{Elm11}. Nevertheless, theoretical IMF models, such as those proposed by \citet{Pad97,Pad02} or \citet{Hen08,Hen09}, provide analytical solutions that help elucidate the contribution of the various physical processes involved, but, to obtain these analytic solutions it is necessary to adopt a series of assumptions that limit their application to specific cases. The predictions of these theories have been compared to the numerical data from simulations. In particular, \citet{Pad04} and \citet{Pad07} compared the analytical solutions of their theory to numerical simulations. \citet{Sch10} have also compared the results from their simulations to the predictions of these theories and have shown how the clump mass distribution depends on the turbulence driving mechanism. In between these two approaches are the phenomenological models, such as the one presented here, that allow one to address some of the questions stated above, in particular, that of sensitivity to the initial conditions. The method consists of a selection criterion based on the thermal and turbulent Jeans mass which is applied hierarchically from small to large scales. \subsection{Aim of the Paper} The methodology proposed in this work enables a direct connection between the structure of molecular clouds and the distributions of generated cores in both mass and space. Thus our first aim is to check that the results obtained using our methodology are consistent with those obtained using other methods proposed in the literature and that have produced reliable results. In particular, we will compare our results for a lognormal density PDF with the theoretical CMF derived by \citet{Hen08}, but using two different spatial distributions of the cloud mass: a) a ramdon cloud, and b) what we have called a ``corner" distribution where the voxel mass decreases with the distance to a preselected corner. This exercise allows us to evaluate the virtue of the method and how the geometry of the cloud defines the dependence of the standard deviation of the lognormal density PDF with the smoothing scale R. We have chosen these two very different spatial structures so as to make it clear how the analytic formulation of HC08 and the phenomenology presented here are connected through the scale dependence of the density PDF. Second, we will explore the formation of cores for different parent clouds, but considering that the geometry that best describes the spatial structure of the clouds is fractal. Observations of close star-forming clouds indicate that the mass distribution in them can be described as having a fractal structure \citep{Fal92,San05}. The analysis will be carried out for fractional Brownian motion (fBm) clouds with a wide range of fractal dimensions. Specifically, we will focus on the comparative analysis of the following properties: a) Empirical dependence of density PDF on the smoothing scale; b) Mass distribution of the cores; c) Spatial distribution of the generated cores as measured by the surface density of companions, and, d) Cloud core properties averaged over several realizations for each $H$. The paper is organized into five sections, this introduction being the first. $\S$2 describes the method and defines the main physical variables of the problem and their range of values in our simulations. In $\S$3 we check the virtue of our methodology in reproducing the CMF derived analytically by Hennebelle \& Chabrier (2008), and in $\S$4, we show the application of this methodology to fractal clouds generated as fBm clouds with different Hurst exponents and compare the results with previous approaches to the same physical systems. Finally, $\S$5 is devoted to summarizing the main conclusions. | We have proposed a procedure to obtain the prestellar core mass distribution that results from the collapse of prestellar cores in clumps by assuming that the densest regions collapse first and form the smaller objects. At small scales, thermal support dominates and determines the mass distribution of cores at low masses, whereas at the largest scales turbulence dominates the support and determines themass distribution at high masses. The numerical method proposed here make use of a small number of parameters, namely $N_{vox}, T, \mu, \eta$, together with the cloud properties (i.e. $M_{cl}$ and $H$) and the assumptions that the mass in the densest voxel $m_{max}$ is equal to the Jeans mass and that the cloud as a whole is marginally stable. When the proposed method is applied to a mass distribution whose density PDF is a log-normal at all smoothing scales $R$ and its standard deviation $\sigma(R)$ is given by eq. (\ref{sig2_R}), the average mass distribution agrees with the CMF predicted by the analytical theory of the IMF proposed by HC08. The HCS method can be seen as a numerical version of the HC08 theory, and there is univocal correspondence between the parameters in both models. Both the Padoan-Nordlund IMF and the Hennebelle-Chabrier IMF apply to particular star-forming cloud conditions. In order to determine an average IMF that can be compared with observations of stars from different clouds, it is necessary to average their theoretical IMFs for a distribution of cloud temperatures, densities and Mach numbers. The core mass distribution from our method is even more dependent on cloud property since, as we have shown, the masses of the resulting cores also depend on the particular distribution of mass within the cloud. Large variations in the resulting core mass distribution are observed in fBm clouds with the same mass $M_{cl}$ and Hurst exponent $H$, but a different setup of the random Fourier phases. As shown in Table 2 and Fig. (\ref{fig:varcore_H}), the number of cores in a set of 20 simulations display variations that largely exceed the expected $\sqrt{\cal N}$ statistical variations. Due to its simplicity the HCS method is computationally efficient at obtaining the mass and position of the cores that collapse in an arbitrary distribution of gas. Therefore the HCS method is well suited to analyzing the effects produced by changes in the physics over a large number of initial conditions. We have applied the HCS method to lattices with a number of cells up to $2^{8\times 3}$, which represent clumps of mass $\sim 10^3 \, \emsun$ and size $\sim 3$ pc, but larger lattices can be processed. There is no restriction in the way the mass in the voxels is assigned, but we have focused on fBm clouds that have been used as analogs of real interstellar clouds. We confirm that fBm clouds with $H\simeq1/3$, corresponding to $\gamma=11/3$ (Elmegreen 2002), give better agreement with the theoretical CMF derived by Hennebelle and Chabrier and the observed IMF. We have also shown that the spatial distribution of the cores for fBm clouds with $H=1/3$ has a surface density of companions that resembles that of young stellar clusters (Simon 1997). Since the HCS method provides the sequence and location of newly formed stars, the method can be easily modified to consider radiative feedback effects. | 12 | 6 | 1206.0873 |
1206 | 1206.2862_arXiv.txt | We present the orbital and physical parameters of a newly discovered low-mass detached eclipsing binary from the \emph{All-Sky Automated Survey} (ASAS) database: {ASAS~J011328-3821.1~A} -- a member of a visual binary system with the secondary component separated by about 1.4 seconds of arc. The radial velocities were calculated from the high-resolution spectra obtained with the 1.9-m Radcliffe/GIRAFFE, 3.9-m AAT/UCLES and 3.0-m Shane/HamSpec telescopes/spectrographs on the basis of the \todcor technique and positions of \ha emission lines. For the analysis we used $V$ and $I$ band photometry obtained with the 1.0-m Elizabeth and robotic 0.41-m PROMPT telescopes, supplemented with the publicly available ASAS light curve of the system. We found that ASAS~J011328-3821.1~A is composed of two late-type dwarfs having masses of $M_1 = 0.612 \pm 0.030$ M$_\odot$, $M_2 = 0.445 \pm 0.019$ M$_\odot$ and radii of $R_1 = 0.596 \pm 0.020$ R$_\odot$, $R_2 = 0.445 \pm 0.024$ R$_\odot$, both show a substantial level of activity, which manifests in strong \ha and $H_\beta$ emission and the presence of cool spots. The influence of the third light on the eclipsing pair properties was also evaluated and the photometric properties of the component B were derived. Comparison with several popular stellar evolution models shows that the system is on its main sequence evolution stage and probably is more metal rich than the Sun. We also found several clues which suggest that the component B itself is a binary composed of two nearly identical $\sim 0.5$ M$_\odot$ stars. | In the last decade there has been a stunning increase in our knowledge of the structure and evolution of eclipsing binaries. The number of known and well studied detached low-mass eclipsing systems, which are the main source of absolute parameter measurements, increased by almost an order of magnitude after the year 2000. This increase helped several authors make conclusions about a 30-year old problem: the discrepancy between observed and theoretically predicted radii and temperatures during the evolution of low-mass stars \citep{lac77}. The prevailing idea is that this discrepancy is caused by the stellar magnetic field, amplified by the fast rotation due to tidal locking in close binaries, which inhibits the efficiency of convection in the envelope and manifests itself in a higher level of activity than for single stars \citep{cha07}. However, the number of known well-characterized systems is too low to reliably test this hypothesis. This is especially true amoung the long-period systems, which normally exhibit lower levels of activity, so the observed radii and temperatures are closer to the theoretically predicted values. In a very recent paper, \citet{irw11} described an eclipsing pair of M dwarfs on a 41 d orbit which, despite slow asynchronous rotation and a probable age over 120 Myr, still show cold spots, and radii are inflated by 5\% with respect to theoretical predictions. This system is evidence against the theory mentioned above and may again bring wider interest to the need to revisit the equation of state for low mass stars, which was suggested to explain the inflated radii before the activity hypothesis \citep[e.g.][]{tor02}. To be useful for testing the evolutionary models, the parameters of the system, especially masses and radii, should be known with a precision of at least 3\%, considered now as a canonical level \citep{bla08,cla08}. There are two other criteria which can make a system more interesting: (1) additional multiplicity -- by deriving properties of the third body one can put additional constraints on the nature of the whole system; (2) low mass ratio -- in general it is harder to fit a single isochrone to data points that are well separated in parameter space; these are usually stars of different masses. Twin stars commonly have data points that lie close together and are less difficult to fit. In this paper we present our analysis together with orbital and physical parameters of a newly discovered low-mass detached eclipsing binary which meets the two mentioned criteria. The data we gathered fell short of the 3\% precision level goal, but the system is still highly valuable for our understanding of the nature of low-mass stars. | \subsection{Evolutionary status} We compare our results with three sets of theoretical isochrones: Yonsei-Yale \citep[Y$^2$;][]{yi01,dem04}, Dartmouth \citep{dot07}, and Padova \citep{gir00,mar08}. \begin{figure} \centering \includegraphics[width=0.79\columnwidth]{evo_01.eps} \caption{Comparison of our results for ASAS-01~A eclipsing pair with theoretical 1 gyr isochrones on mass $M$ vs. (from top to bottom): bolometric magnitude, radius, effective temperature, absolute $V$ magnitude, and $V-I$ colour. \yy isochrones are depicted with solid, Darthmouth with dashed, and Padova with dot-dashed line. Isochrones for solar metallicity ($Z\simeq 0.02$) are depicted with red and for higher $Z$ with green. The colour version of the picture is available in the on-line version of the paper.}\label{fig_evo} \end{figure} The comparison is shown in Figure \ref{fig_evo}. We present the bolometric magnitude, radius, effective temperature, absolute $V$ magnitude and the $V-I$ colour as a function of stellar mass. We compare our results with isochrones for an age of 1 Gyr and two cases of metallicities: (1) solar, $Z\simeq0.02$ (red lines), and (2) above solar (green lines), however different for every set: $Z=0.04$ for Y$^2$, 0.035 for Dartmouth, and 0.03 for Padova. Since the accuracy of our measurements is worse than the canonical 3\%, the isochrones are shown only for comparison, i.e. we do not attempt to determine the age of the system. It is however enough to conclude that ASAS-01 is probably a main sequence object. The metallicity determination is even more insecure, however from the Fig. \ref{fig_evo} one can claim that metallicities higher than solar are prefered -- on the $M$ vs. $M_{bol}$ and $\log{T_{eff}}$ planes the $Z>0.02$ isochrones fit significantly better to the data points. The main sequence evolutionary stage of ASAS-01 is confirmed by its galactic kinematics. We used our determinations of the systemic velocity and distance together with the position and proper motion from the PPMXL catalogue: $\mu_\alpha = 120.4\pm4.0$~mas~yr$^{-1}$, $\mu_\delta = -36.7\pm4.0$~mas~yr$^{-1}$ \citep{roe10}. The obtained values of $U = 3.3\pm2.3$~km~s$^{-1}$, $V = -16.8\pm2.1$~km~s$^{-1}$ and $W = -7.8\pm2.3$~km~s$^{-1}$ put ASAS-01 in the galactic thin disk, and only marginally in one of the young moving groups recently reported by \citet[][ID=15 therein]{zha09}. It is worth noticing that on the $M/R$ plane the measurements agree within errors with the \yy and Dartmouth models. The second one especially seems to reproduce the physical properties of ASAS-01~A components. It is in a somewhat contrary to many recent results of low-mass detached eclipsing binaries studies \citep[see:][]{kra11}. The usually, but not always, observed characteristic of low-mass stars in close binaries is that theoretically predicted radii are smaller and temperatures larger with respect to what is observed \citep{rib08}, and this situation is seen for the Padova models. The Dartmouth models also seem to follow that trend, but only by a few per cent. Within our uncertainties they correctly predict the observed radii, and, in the case of $Z=0.035$, also temperatures. The deviation is smaller for the larger component, which presumably rotates faster, thus it does not support the possible radius-rotation relation for $M<0.7$~M$_\odot$ stars \citep{kra11}. At the same time the properties of the secondary are very well reproduced by the \yy models (except for the colour) but for the primary we see the ``typical'' underestimation of the radius and overestimation of the temperature. From this short discussion we can only conclude that the Dartmouth models may be the best to reproduce the properties of low-mass stars \citep[see also:][]{tho10} although discrepancies at the level of $\sim$3\% are present. This is however significantly smaller than for the majority of previous studies where 5-15\% difference in radii were claimed. \subsection{The companion B}\label{sec_B} Having the two components separated on $V$ and $I$ band images (Fig. \ref{fig_foto}), and the physical model created with \phoebe with the third light included, we were able to estimate the photometric properties of component B. We present the fractional fluxes in both bands, the dereddened $V-I$ colour, and the absolute $V$ magnitude in Table \ref{tab_3rd}. \begin{table} \centering \caption{Photometric properties of the component B from the \phoebe model}\label{tab_3rd} \begin{tabular}{lrl} \hline \hline Parameter & Value & $\pm$ \\ \hline $F_{V,3}$ [\%]$^a$ & 34.8 & 1.5 \\ $F_{I,3}$ [\%]$^a$ & 42.9 & 3.2 \\ $(V-I)_3$ [mag]$^b$ & 2.15 & 0.13 \\ $M_{V,3}$ [mag] & 10.82 & 0.70 \\ \hline \end{tabular} \flushleft $^a$ Fractional fluxes defined as a percentage of the flux of component A\\ $^b$ Colour dereddened with $E(V-I)=0.02$ mag \end{table} \begin{table} \centering \caption{Values of the inclination of the component B putative orbit for various values of $q_B$ and masses of the more massive component of 0.45 M$_\odot$ ($i_{B,0.45}$) and 0.5 M$_\odot$ ($i_{B,0.5}$). The mass function is taken from the Table \ref{tab_rv_B}.}\label{tab_q_3rd} \begin{tabular}{ccccc} \hline \hline $q_B$ & $M_{Ba}\sin^3{i_B}$ & $M_{Bb}\sin^3{i_B}$ & $i_{B,0.45}$ & $i_{B,0.5}$\\ & [M$_\odot$] & [M$_\odot$] & [$^\circ$] & [$^\circ$] \\ \hline 1.00 &0.226 &0.226 &53 &50\\ 0.95 &0.251 &0.238 &55 &53\\ 0.90 &0.280 &0.252 &59 &56\\ 0.85 &0.315 &0.268 &63 &59\\ 0.80 &0.358 &0.286 &68 &63\\ ... &&&& \\ 0.718 &0.45 &0.323 &90 &75\\ 0.684 &0.5 &0.342 &--- &90\\ \hline \multicolumn{5}{l}{$f(M_{Ba}) = 0.0565(92)$ M$_\odot$} \\ \end{tabular} \end{table} \begin{figure} \centering \includegraphics[width=\columnwidth]{3rdlight.eps} \caption{Colour-brightness diagram with our measurements for ASAS-01~Aa, Ab and B components, together with theoretical 1~Gyr isochrones. Line and colour coding is the same as in Fig. \ref{fig_evo}. The square shows the approximate position of the components of ASAS-01~B if it was composed of two identical stars showing the same $V-I$ value as observed. The colour version of the Figure is available in the on-line version of the manuscript.}\label{fig_3rd} \end{figure} In Figure \ref{fig_3rd} we present our measurements on the $V~-~I/M_V$ plane. Data from Tables \ref{tab_par} and \ref{tab_3rd}, and the same isochrones as in Fig. \ref{fig_evo} are used. We see that again the Dartmouth set for $Z=0.02$ gives the best match for the eclipsing pair A. Component B is about 0.7~mag brighter but has almost the same colour as Ab ($V-I$ = 2.14 mag). Due to relatively large errorbars (magnified mainly by the uncertainty in the third light parameter in our model) one can formally find $V$ magnitudes of B and Ab almost equal and still rather consistent with the considered Dartmouth and both \yy isochrones. However, if assuming that component B is composed of two nearly-twin stars with masses $\sim$0.45-0.5~M$_\odot$, which would thus have the $V-I$ colour around 2.15~mag, i.e. almost the same as the value observed for B, those putative Ba and Bb components would be located in Fig. \ref{fig_3rd} on the position marked by the empty square (lowering the flux by a factor of 2 increases the magnitude by $\sim$0.7). In such a situation, the Darthmouth $Z=0.02$ and \yy isochrones would reproduce the observed properties of the system much better. We consider this fact as a support for the hypothesis of the binarity of ASAS-01~B. The binarity of component B, with a mass ratio close to 1 and component colours close to the value observed, is consistent with the fact that no marks of additional eclipses are found in our $V$ and $I$ band light curves. From the orbital fit for component B radial velocities, presented in Table \ref{tab_rv_B}, we can estimate the orbital inclination $i_B$ for various values of the mass-ratio $q_B$ and the desired mass of the more massive component $M_{Ba}$. We present those calculations in Table \ref{tab_q_3rd}. For a number of different values of $q_B$, we present values of $M \sin^3{i_B}$ for Ba and Bb, and values of inclination angles, for which $M_{Ba} = 0.45$ or 0.5 M$_\odot$. From this Table one can see, that if ASAS-01~B is a binary, to explain all the observed properties one needs the inclination angle $i_B$ to be between 50 and 70 degrees, which corresponds to $q\ge 0.8$. For these values of masses of Ba and Bb we can expect the observed colour to be around the value given in Tab. \ref{tab_3rd}. Assuming lower values of $q_B$ we would end up in a situation where additional eclipses occur and the Bb component is significantly fainter. In such a case the total brightness of ASAS-01~B would be lower than it is observed, however still within relatively large error bars. If ASAS-01 is a double-double system, it would join a small group of this kind of interesting objects. To date only two other low-mass detached eclipsing binaries are known to be in similar configuration: BD~-22~5866 -- a system with K7+K7 eclipsing pair and M1+M2 non-eclipsing binary \citep{shk08}, and YY~Gem -- the faintest member of a sextuple system composed of three spectroscopic binaries -- Castor A, B and C \citep[$\alpha$~Gem~ABC;][]{vin40,kro52,bop74}. There are also few examples of LMDEBs known to have a single additional companion, like LP~133-373 \citep{vac07}, HIP~96515\footnote{No RV curve is published for this system.} \citep{hue09}, MR~Del \citep{pri09,dju11}, NLTT~41135 \citep{irw10}, ASAS-08 \citep{mon07,hel11a} or the triply-eclipsing KOI-126 \citep{car11}. Such systems not only allow more rigid constrains on the evolutionary models than in the cases of ``lonely'' eclipsing binaries, but also play important role in testing star-formation theories, stellar population codes and dynamical interactions in multiple stellar systems. The relative brightness and small distance to ASAS-01 make it a valuable object for further studies. | 12 | 6 | 1206.2862 |
1206 | 1206.2115_arXiv.txt | Young stars form in molecular cores, which are dense condensations within molecular clouds. We have searched for molecular cores traced by $^{13}$CO $J=1\to 0$ emission in the Taurus molecular cloud and studied their properties. Our data set has a spatial dynamic range (the ratio of linear map size to the pixel size) of about 1000 and spectrally resolved velocity information, which together allow a systematic examination of the distribution and dynamic state of $^{13}$CO cores in a large contiguous region. We use empirical fit to the CO and CO$_2$ ice to correct for depletion of gas-phase CO. The $^{13}$CO core mass function ($^{13}$CO CMF) can be fitted better with a log-normal function than with a power law function. We also extract cores and calculate the $^{13}$CO CMF based on the integrated intensity of $^{13}$CO and the CMF from 2MASS. We demonstrate that there exists core blending, i.e.\ combined structures that are incoherent in velocity but continuous in column density. The core velocity dispersion (CVD), which is the variance of the core velocity difference $\delta v$, exhibits a power-law behavior as a function of the apparent separation $L$:\ CVD (km/s) $\propto L ({\rm pc})^{0.7}$. This is similar to Larson's law for the velocity dispersion of the gas. The peak velocities of \13co\ cores do not deviate from the centroid velocities of the ambient \co\ gas by more than half of the line width. The low velocity dispersion among cores, the close similarity between CVD and Larson's law, and the small separation between core centroid velocities and the ambient gas all suggest that molecular cores condense out of the diffuse gas without additional energy from star formation or significant impact from converging flows. | Most young stars are found in dense molecular cores~\citep{McKee2007}. There is a large volume of data concerning dense molecular cores traced by dust emission and dust extinction~\citep{Motte1998,Testi1998,Johnstone2000,Stanke2006,Reid2005,Reid2006,Johnstone2001,Johnstone2006,DCMF_Alves_Lombardi_Lada}. One of the primary outcomes of these studies is the core mass function (CMF). Because of the still unknown physical origin of the stellar initial mass function (IMF) and its significance, emphasis has been placed on the possible connection between the CMF and the IMF~\citep{DCMF_Alves_Lombardi_Lada}. The construction of an analytic form of the CMF from observational data has largely focused on two functional forms, power law and log-normal. The majority of past studies claim to find a power law CMF, the shape of which resembles the Salpeter IMF~\citep{IMF} (The upper mass limit of the original Salpeter IMF is $10 M_{\odot}$) \begin{equation} \frac{{\rm d}N}{{\rm d}\log M} \propto M^{-\gamma},\ \ \gamma=1.35, -0.4\le\log(M/M_{\odot})\le 1.0 \lp \label{power_law} \end{equation} Such a power law CMF was found in millimeter continuum maps of the $\rho$ Ophiuchus region by \cite{Motte1998} with the IRAM 30-meter telescope and a series of subsequent studies in the Serpens region \citep{Testi1998}, the $\rho$ Ophiuchus region \citep{Johnstone2000,Stanke2006}, NGC 7538 \citep{Reid2005}, M17 \citep{Reid2006}, Orion \citep{Johnstone2001,Johnstone2006} and the Pipe nebula \citep{DCMF_Alves_Lombardi_Lada}. \cite{Reid2006b} studied the CMF in 11 star-forming regions and find an average power law index of $\gamma=1.4\pm 0.1$. Recently, dust emission observations of the Aquila rift complex with {\it Herschel} \footnote{http://www.esa.int/SPECIALS/Herschel/index.html} reveal a power law mass function with $\gamma=1.5\pm0.2$, for $M>2\ M_{\odot}$~\citep{Aquila_cores}, which is also consistent with the Salpeter IMF. A similar power law index is found in some molecular emission studies. For example, the CMF obtained from a $\rm C^{18}O$ study in the S140 region has $\gamma=1.1\pm 0.2$~\citep{S140}. On the other hand, there are also studies that find a flatter CMF. \cite{Kramer1998} studied seven molecular clouds L1457, MCLD 126.6+24.5, NGC 1499 SW, Orion B South, S140, M17 SW, and NGC 7538 in $^{13}$CO and C$^{18}$O, and find $\gamma$ to be between 0.6 and 0.8. \cite{core_mass_function} studied cores in the Orion molecular cloud traced by sub-millimeter continuum and found a power law with $\gamma=0.15\pm 0.21$. These findings of small $\gamma$ are in the minority and do not seem to be a special result of spectroscopic mapping. \cite{Pipe_nebula_improve} obtained the CMF from an extinction map and used complementary C$^{18}$O observations to examine the effects of blending of cores in dust maps. They claimed to find a CMF with $\gamma$ similar to that of the Salpeter IMF. At a first glance, a similarity between the CMF and Salpeter IMF suggests a constant star formation efficiency, which is independent of the core mass. It is crucial to note, however, these studies are examining structures on vastly different scales of size, mass, and density. In \cite{Reid2006}, for example, the mass of the cores ranges from about 0.1 $M_{\odot}$ to $1.6\times 10^4\ M_{\odot}$. Due to the large distances of many targeted regions, any "cores" over about 500 \Ms\ are certainly unresolved, with many showing signs of much evolved star formation, such as water masers \citep{Wang2006} and/or compact HII regions \citep{Hofner2002}. The observed similarity between CMF and Salpeter IMF may be explained equally well by self-similar cloud structures as well as a constant star formation efficiency. Some observations (e.g., \cite{observe_lognormal}) suggest a log-normal form for the CMF (in the mass range of $0.1\ M_{\odot}<M<10\ M_{\odot}$), \begin{equation} \frac{{\rm d}N}{{\rm d}\log M} \propto \exp\left[-\frac{(\log M-\mu)^2}{2\sigma^2}\right]\lp \label{lognormal} \end{equation} This is reminiscent of the \cite{Chabrier2003} IMF, which is also of log-normal form. Theoretically, if the core mass depends on $n$ quantities which are random variables, the CMF would be log-normal when $n$ is large, i.e., the core formation processes are complicated~\citep{adams1996}. This is a result of the central limit theorem. A log-normal distribution also arises naturally from isothermal turbulence~\citep{Larson1973}. It is thus of great interest to distinguish the two forms of the CMF and obtain the key parameters associated with each form. \cite{Form_of_CMF} show that a large sample with many cores is needed to differentiate these two forms. Furthermore, we also emphasize here the critical need to obtain a large sample of cores in spectroscopic data. Overlapping cores along the same line of sight can only be separated using resolved velocity information. It is important to evaluate the effect of such accidental alignment on the derived CMF. A Nyquist sampled continuous spectroscopic map is also essential for studying the dynamical characteristics of star forming regions, such as the Core Velocity Dispersion (CVD $\equiv \langle\delta v^2\rangle^{1/2}$ see section 4.3), of the whole core sample in one star forming region. The Taurus molecular cloud is a nearby (with a distance of 140 pc,~\cite{Distance}) low-mass star-forming region. In this work, we obtain a sample of cores in the $^{13}$CO data cube of this region. We study the properties of $^{13}$CO cores in detail and compare them with those found in the dust extinction map of the same region. We first briefly describe the data in \S \ref{sec:data}; we present the methods used to find $^{13}$CO cores in \S \ref{sec:methods}; we present the observed $^{13}$CO CMF and CVD in \S \ref{sec:results}; we discuss the implications of our observations in \S \ref{sec:discussion}. In the final section we present our conclusions. | \label{sec:conclusion} We have studied the $^{13}$CO cores identified within the Taurus molecular cloud using a 100 $\rm degree^2$ \13co\ $J=1\to 0$ map of this region. The spatial resolution of 0.014 pc and the velocity resolution of 0.266 km/s facilitate a detailed study of the physical conditions of $^{13}$CO cores in Taurus. The spatial dynamic range (the ratio of linear map size to the Nyquist sampling interval) of 1000 of our data set allows examination of the collective motions of $^{13}$CO cores and their relationship to their surroundings. We have found that the velocity information helps to exclude cores which we consider to be spurious. Our conclusions regarding the extraction of $^{13}$CO cores and their properties are: 1) Velocity information is essential in resolving overlapping cores, allowing better resolution of cores and a more accurate determination of the CMF. 2) The mass function of the 3D(x,y,v) \13co\ cores can be fitted better with a log-normal function ($\mu = 0.95$ and $\sigma= 1.15$) than with a power law function. For cores represented by \13co\ total intensity and extinction, a log-normal distribution is also a better representation of the mass distribution than is a power law. There is no simple relation between the Taurus $^{13}$CO CMF and the stellar IMF. 3) No $^{13}$CO cores are found to have mass greater than the critical Bonnor-Ebert mass, in contrast to cores in Orion. 4) Only 10\% of $^{13}$CO cores are approximately bound by their own gravity (with the virial parameter $M_{\rm vir}/M<2$). These can be well fitted with a log-normal mass function. External pressure plausibly plays a significant role in confining the $^{13}$CO cores with small density contrast to the surrounding medium. 5) In Taurus, the relation between core velocity dispersion (CVD$\equiv \langle\delta v^2\rangle^{1/2}$) and the apparent separation between cores $L$ can be fitted with a power law of the form CVD (km/s) $=0.2L ({\rm pc})^{0.7}+0.2$ in the $ 0 \leq L \leq 10$ pc region, similar to the Larson's law, with a median value of 0.78 km/s. 6) The observed CVD is reproduced by using a simple two-core-cluster model, in which there are two core clusters with radius 7 pc and a separation 9 pc between these two clusters. 7) The low velocity dispersion among cores, the close similarity between CVD and Larson's law, and the small difference between core centroid velocities and the ambient diffuse gas all suggest that dense cores condense out of the diffuse gas without additional energy input and are consistent with an ISM evolution picture without significant feedback from star formation or significant impact from converging flows. 8) The CVD can be an important diagnostic of the core dynamics and the cloud evolution. We encourage simulators to provide comparable information based on their calculations. | 12 | 6 | 1206.2115 |
1206 | 1206.6500_arXiv.txt | We have observed 37 Infrared Dark Clouds (IRDCs), containing a total of 159 clumps, in high-density molecular tracers at 3 mm using the 22-meter ATNF Mopra Telescope located in Australia. After determining kinematic distances, we eliminated clumps that are not located in IRDCs and clumps with a separation between them of less than one Mopra beam. Our final sample consists of 92 IRDC clumps. The most commonly detected molecular lines are (detection rates higher than 8\%): N$_2$H$^+$, HNC, HN$^{13}$C, HCO$^+$, H$^{13}$CO$^+$, HCN, C$_2$H, HC$_3$N, HNCO, and SiO. We investigate the behavior of the different molecular tracers and look for chemical variations as a function of an evolutionary sequence based on {\it Spitzer} IRAC and MIPS emission. We find that the molecular tracers behave differently through the evolutionary sequence and some of them can be used to yield useful relative age information. The presence of HNC and N$_2$H$^+$ lines do not depend on the star formation activity. On the other hand, HC$_3$N, HNCO, and SiO are predominantly detected in later stages of evolution. Optical depth calculations show that in IRDC clumps the N$_2$H$^+$ line is optically thin, the C$_2$H line is moderately optically thick, and HNC and HCO$^+$ are optically thick. The HCN hyperfine transitions are blended, and, in addition, show self-absorbed line profiles and extended wing emission. These factors combined prevent the use of HCN hyperfine transitions for the calculation of physical parameters. Total column densities of the different molecules, except C$_2$H, increase with the evolutionary stage of the clumps. Molecular abundances increase with the evolutionary stage for N$_2$H$^+$ and HCO$^+$. The N$_2$H$^+$/HCO$^+$ and N$_2$H$^+$/HNC abudance ratios act as chemical clocks, increasing with the evolution of the clumps. | Although far less common than low-mass stars, massive stars play a key role in the evolution of the energetics and chemistry of molecular clouds and galaxies. However, the formation of high-mass stars is far less clear than their low-mass counterparts for several reasons. Massive stars are rare and evolve quickly. In addition, the regions that host their early stages of formation are difficult to observe due to high dust extinction, and, with a few exceptions, are located at large distances ($\gtrsim$ 3 kpc). For this reason, most of the research on massive star formation has been based on observations of regions with current star formation (e.g., ``hot cores,'' \hii\ regions), which were initially easier to detect in surveys and to follow up in detail. In constrast, objects in the earlier ``prestellar'' or ``starless'' stage have been much harder to find and, in consequence, this stage still remains poorly understood. About a decade ago, however, the regions containing the earliest stages of massive star formation were identified. Galactic plane surveys revealed thousands of dark patches obscuring the bright mid-infrared background ({\it ISO}, \citealt{Perault96}; {\it MSX}, \citealt{Egan98}, \citealt{Simon06a}; {\it Spitzer}, \citealt{Peretto09}, \citealt{Kim10}). These dark silhouettes were associated with molecular and dust emission, indicating they consist of dense molecular gas. Such objects were called Infrared Dark Clouds (IRDCs). The first studies characterizing them suggested that they were cold ($<$25 K), massive ($\sim$10$^2$--10$^4$ \Msun), and dense ($\gtrsim$10$^5$ cm$^{-3}$) molecular clouds with high column densities ($\gtrsim$10$^{23}$ cm$^{-2}$) \citep{Carey98,Carey00}. They also correspond to the densest clouds embedded within giant molecular clouds \citep{Simon06b}. More recent studies on IRDC clumps,\footnote{Throughout this paper, we use the term ``clump'' to refer to a dense object within an IRDC with a size of the order $\sim$1 pc and a mass $\sim$10$^2$--10$^3$ \Msun. We use the term ``core'' to describe a compact, dense object within a clump with a size $\lesssim$0.1 pc and a mass $\lesssim$50 \Msun.} using dust continuum emission, show that they have typical masses of $\sim$120 \Msun\ and sizes $\sim$0.5 pc \citep{Rathborne06}. Spectral energy distributions (SEDs) of several IRDC clumps reveal dust temperatures that range between 16 and 52 K, and luminosities that range from $\sim$10--10$^{5}$ \Lsun\ \citep{Rathborne10}. Temperatures derived using NH$_3$ observations are lower than dust temperatures and range from $\sim$10--20 K \citep{Pillai06,Sakai08,Devine11,Ragan11}. Star formation activity in IRDC clumps can be inferred from high temperatures and high luminosities as well as the presence of UC \hii\ regions \citep{Battersby10}, hot cores \citep{Rathborne08}, embedded 24 $\mu$m sources \citep{Chambers09}, molecular outflows \citep{Beuther07,Sanhueza10}, and maser emission \citep{Wang06,Chambers09}. The high masses, densities, and column densities of the IRDC clumps, as well as the aforementioned signatures of star formation in them, indicate that IRDCs are currently forming massive stars. In addition, IRDCs harbor numerous candidates for the most elusive earliest phase of high-mass star formation, the ``prestellar'' or ``starless'' phase \citep{Chambers09,Rathborne10,Rygl10,Vasyunina11,Pillai11,Devine11}. Due to the characteristics of IRDCs and the large variety of evolutionary stages that they harbor, it has been suggested that they are the natal sites of all massive stars and stellar clusters \citep{Rathborne06,Rathborne10}. With the aim of characterizing the different evolutionary stages of clumps found in IRDCs, \cite{Chambers09} proposed an evolutionary sequence in which ``quiescent'' clumps evolve into ``intermediate'', ``active'', and ``red'' clumps. This evolutionary scheme is based on the {\it Spitzer}/IRAC 3-8 $\mu$m colors and the presence or absence of {\it Spitzer}/MIPS 24 $\mu$m point-source emission. A clump is called ``quiescent'' if it contains no IR-{\it Spitzer} emission (it is IR-dark); ``intermediate'' if it contains either an enhanced 4.5 $\mu$m emission, the so-called ``green fuzzies'' \citep[also known as Extended Green Objects (EGOs);][]{Cyganowski08}, or a 24 $\mu$m source, but not both; ``active'' if it is associated with a green fuzzy and an embedded 24 $\mu$m source; and ``red'' if it is associated with bright 8 $\mu$m emission, which likely corresponds to an \hii\ region. Quiescent clumps are the best candidates to be in the ``prestellar'' or ``starless'' phase, and the massive quiescent clumps are the places where it is most probable that high-mass stars are in their earliest stages of evolution. An additional category is ``blue,'' which describes objects with bright 3.6 $\mu$m emission that are predominantly unextincted stars. Although several properties of IRDCs have been determined in the last few years, there is still one that remains poorly explored, their chemistry. What is the chemistry in IRDC clumps? Are the evolutionary stages defined by \cite{Chambers09} chemically distinguishable? Presently, only a few studies have focused on this subject \citep{Sakai08,Sakai10,Sakai12,Vasyunina11,Miettinen11,Chen11}. In this paper, we report initial results from a program aimed at better understanding the chemical evolution of IRDC clumps, carrying out a multi-line survey at 3 mm. We observed several molecular lines simultaneously, which facilitates comparison between different lines by eliminating or reducing observational errors arising from uncertainties in telescope pointing and calibrations. The main goals of this project are to investigate the behavior of the different molecular tracers and look for observable changes in the chemistry as a function of the evolutionary stages proposed by \cite{Chambers09}. \begin{deluxetable*}{lcccccc} \tabletypesize{\scriptsize} \tablecaption{Summary of Observed Molecular Lines \label{tbl-obsparam}} \tablewidth{0pt} \tablehead{ \colhead{Molecule} & \colhead{Transition} & \colhead{Rest Frequency} & \colhead{$E_u/k$} & \colhead{$n_{\rm crit}$}& \colhead{IF} & \colhead{$T_{\rm rms}$} \\ \colhead{} & \colhead{} & \colhead{(GHz)}&\colhead{(K)}&\colhead{(cm$^{-3}$)}&\colhead{}&\colhead{(K)}\\ } \startdata \NHdosD & \tNHdosD & 85.926260 & 20.68& $4\times 10^6$ & IF3 & 0.044 \\ \SO & \tSO & 86.093983 & 19.31& $2\times 10^5$ & IF3 & 0.044 \\ \HtreceCN & \tHtreceCNuu& 86.338735 & 4.14 & $2\times 10^6$ & IF3 & 0.044 \\ & \tHtreceCNdu& 86.340167 & 4.14 & $2\times 10^6$ & IF3 & 0.044 \\ & \tHtreceCNcu& 86.342256 & 4.14 & $2\times 10^6$ & IF3 & 0.044 \\ \HtreceCO & \tHtreceCO & 86.754330 & 4.16 & $2\times 10^5$ & IF3 & 0.044 \\ \SiO & \tSiO & 86.846998 & 6.25 & $2\times 10^6$ & IF3 & 0.044 \\ \HNtreceC & \tHNtreceC & 87.090859 & 4.18 & $3\times 10^5$ & IF3 & 0.044 \\ \CdosH & \tCdosHudcu & 87.316925 & 4.19 & $2\times 10^5$\tablenotemark{a} & IF3 & 0.044 \\ & \tCdosHuucc & 87.328624 & 4.19 & $2\times 10^5$\tablenotemark{a}& IF3 & 0.044 \\ & \tCdosHuucu & 87.402004 & 4.19 & $2\times 10^5$\tablenotemark{a} & IF3 & 0.044 \\ & \tCdosHuccu & 87.407165 & 4.19 & $2\times 10^5$\tablenotemark{a} & IF3 & 0.044 \\ \HNCO & \tHNCO & 87.925252 & 10.55& $1\times 10^6$ & IF2 & 0.048 \\ \HCN & \tHCNuu & 88.630416 & 4.25 & $3\times 10^6$ & IF2 & 0.048 \\ & \tHCNud & 88.631847 & 4.25 & $3\times 10^6$ & IF2 & 0.048 \\ & \tHCNuc & 88.633936 & 4.25 & $3\times 10^6$ & IF2 & 0.048 \\ \HCO & \tHCO & 89.188526 & 4.28 & $2\times 10^5$ & IF2 & 0.048 \\ \HNC & \tHNC & 90.663574 & 4.35 & $3\times 10^5$ & IF1 & 0.042 \\ \HCtresN & \tHCtresN & 90.978989 & 24.01& $5\times 10^5$ & IF1 & 0.042 \\ \CHtresCN& \tCHtresCNuu & 91.985316 & 20.39& $4\times 10^5$ & IF0 & 0.042 \\ & \tCHtresCNcc& 91.987089 & 13.24& $5\times 10^5$ & IF0 & 0.042 \\ \CS & \tCS & 92.494303 & 6.00 & $3\times 10^5$ & IF0 & 0.042 \\ \NdosH & \tNdosHuud & 93.171913 & 4.47 & $3\times 10^5$ & IF0 & 0.042 \\ & \tNdosHudt & 93.173772 & 4.47 & $3\times 10^5$ & IF0 & 0.042 \\ & \tNdosHucu & 93.176261 & 4.47 & $3\times 10^5$ & IF0 & 0.042 \\ \enddata \tablecomments{The critical density was calculated as $n_{\rm crit}=A_{\rm ul}/\gamma_{\rm ul}$, where $A_{\rm ul}$ is the Einstein coefficient and $\gamma_{\rm ul}$ is the collisional rate. Values of $A_{\rm ul}$ and $\gamma_{\rm ul}$ at 20 K (50 K for SO) were obtained for most of the molecules from the Leiden Atomic and Molecular Database (LAMDA) \citep{Schoier05}. Values of $A_{\rm ul}$ and $\gamma_{\rm ul}$ at 25 K for \NHdosD\ were obtained from \cite{Machin06}.} \tablenotetext{a}{Critical density adopted from \cite{Lo09}.} \end{deluxetable*} | We have carried out a multi-line survey at 3 mm toward 37 IRDCs, containing 159 clumps, in order to investigate the behavior of the different molecular tracers and search for chemical variations through an evolutionary sequence based on {\it Spitzer} IRAC and MIPS emission. We observed N$_2$H$^+$, HNC, HN$^{13}$C, HCO$^+$, H$^{13}$CO$^+$, HCN, C$_2$H, HC$_3$N, HNCO, and SiO lines with the Mopra 22 m telescope located in Australia. After eliminating clumps that are not located in IRDCs and pairs of clumps that are placed within one Mopra beam, we base our study on 92 sources. HNC and N$_2$H$^+$ lines are detected in almost every IRDC clumps at every evolutionary stage, indicating that their presence does not depend on the star formation activity. On the other hand, HC$_3$N, HNCO, and SiO lines are predominantly detected in later stages of evolution, as expected from their formation paths. The line widths of N$_2$H$^+$ slightly increase with the evolution of the clumps, which is likely produced by the rise of turbulence due to the enhancement of the star formation activity at later evolutionary stages. The increase is modest because, due to the large Mopra beam, we are also tracing the bulk motions of the gas, instead of just the densest regions associated with star formation. Optical depth calculations show that the N$_2$H$^+$ line is mostly optically thin (median of 0.8) and the C$_2$H line is moderately optically thick (median of 2.4). HCO$^+$ and HNC lines are optically thick (medians of 21 and 19, respectively), while their isotopologues are optically thin (median of 0.4 for both). N$_2$H$^+$ opacities show no variations with the evolution of the clumps, whereas C$_2$H, HCO$^+$ and HNC show a slight decrease with the rise of star formation activity. In general, column densities of the different molecules change for the different evolutionary stages defined by \cite{Chambers09} and increase with the evolution of the clumps, with the exception of C$_2$H. However, this is not generally true for molecular abundances (i.e., after dividing by the total H$_2$ column density inferred from 1.2 mm continuum emission). Only the increases of N$_2$H$^+$ and HCO$^+$ abundances are statistically significant and reflect chemical evolution. This is consistent with the results of \cite{Busquet11}, who included both molecules in their chemical modeling of a massive star-forming region. Although it is expected a rise of the HCO$^+$ abundance with the evolution of the clumps, it is not clear why N$_2$H$^+$ also follows this trend. The N$_2$H$^{+}$/HCO$^+$ abundance ratio acts as a chemical clock, increasing its value from intermediate to active and red clumps. This observed trend is consistent with the theoretical predictions. The chemical models suggest that when clumps warm up, they release CO from grain mantles to the gas phase. This rise of CO increases the amount of HCO$^+$, because CO is its main supplier, with respect to N$_2$H$^+$, because CO is its main destroyer. However, the observed trend does not extend to quiescent clumps. This could be due to observational limitations or because the theoretical predictions are incorrect at very early stages of evolution. We also find that the N$_2$H$^{+}$/HNC abundance ratio increases with the evolution of the clumps, from quiescent to red clumps. It is not clear why this ratio behaves in this way, but it suggests that HNC may be preferentially formed in cold gas. | 12 | 6 | 1206.6500 |
1206 | 1206.4219_arXiv.txt | We plan to install an infrared telescope at the new site of Tibet, China. The primary mirror diameter is 50$cm$, and the focal ratio F8. The Xenics 640$\times$512 near infrared camera is employed, equipped with a dedicated high speed InGaAs detector array, working up to 1.7$\mu m$. The new site is located on 5100$m$ mountain, near Gar town, Ali, where is an excellent site for both infrared and submillimeter observations. The telescope will be remotely controlled through internet. The goal of IRT is to make site testing, detect variable stars, and search for extrasolar planets. | The telescope made by Meade, a world leader in manufacturing of amateur telescopes. Its primary mirror diameter is 50$cm$, and the focal ratio is F8. A well sited 50$cm$ telescope could reach the 19 magnitude with a Deep Sky Imager(DSI) by 1 minute. \begin{figure}[h] \begin{minipage}{14pc}\hspace{1pc} \includegraphics[width=13pc]{f1.eps} \caption{The location of Ali site} \end{minipage}\hspace{2pc}% \begin{minipage}{16pc}\hspace{0.5pc} \includegraphics[width=8pc,angle=270]{f2.eps} \caption{The layout of remote control} \end{minipage} \end{figure} The main instruments include a 4K$\times$4K optical camera, and a 640$\times$512 near infrared camera. The STX-16803 optical camera is made by SBIG. The camera has a full frame image buffer for storing image data during download. The XEVA-FPA-1.7-640 infrared camera is one of the revolution in short wave IR cameras, working up to 1.7$\mu m$. The candidate site selected is called Ali, Tibet(Fig1), located at N32$^\circ$19$\verb|'|$, E80$^\circ$01$\verb|'|$, with altitude of 5100$m$\cite{yao08,liu08}. Remote study with the long-term database of ground weather stations and archival satellite data has been performed\cite{qian11,zhang10}. The site has enough relative height on the plateau and is accessible by car. | 12 | 6 | 1206.4219 |
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1206 | 1206.1861_arXiv.txt | It is commonly believed that the earliest stages of star--formation in the Universe were self-regulated by global radiation backgrounds -- either by the ultraviolet Lyman--Werner (LW) photons emitted by the first stars (directly photodissociating ${\rm H_2}$), or by the X-rays produced by accretion onto the black hole (BH) remnants of these stars (heating the gas but catalyzing ${\rm H_2}$ formation). Recent studies have suggested that a significant fraction of the first stars may have had low masses (a few ${\rm M_\odot}$). Such stars do not leave BH remnants and they have softer spectra, with copious infrared (IR) radiation at photon energies $\sim 1$eV. Similar to LW and X--ray photons, these photons have a mean--free path comparable to the Hubble distance, building up an early IR background. Here we show that if soft--spectrum stars, with masses of a few ${\rm M_\odot}$, contributed $\gsim 1\%$ of the UV background (or their mass fraction exceeded $\sim 90\%$), then their IR radiation dominated radiative feedback in the early Universe. The feedback is different from the UV feedback from high-mass stars, and occurs through the photo-detachment of ${\rm H^-}$ ions, necessary for efficient ${\rm H_2}$ formation. Nevertheless, we find that the baryon fraction which must be incorporated into low--mass stars in order to suppress ${\rm H_2}$--cooling is only a factor of few higher than for high-mass stars. | In hierarchical models of structure formation, the first stars in the Universe form in dark matter (DM) minihalos with masses of $\sim 10^5~\msun$ at redshifts of $z\sim 20-30$, through efficient cooling of the gas by ${\rm H_2}$ (\citealt{HTL96}; see a comprehensive review by \citealt{BLreview}). However, soon after the first stars appear, early radiation backgrounds begin to build up, resulting in feedback on star--formation. In particular, the UV radiation in the Lyman--Werner (LW) bands of ${\rm H_2}$ can photodissociate these molecules and suppress gas cooling, possibly preventing star-formation \citep{HRL97,ON99,HAR00,CFA00,MBA01,Ricotti+01, Ricotti+02b,MBH06,WA07,ON08,JGB08,WA08a,WA08b,Whalen+08,MBH09}. Numerical simulations (e.g. \citealt{ABN02,BCL02,Yoshida+03}) have long suggested that the metal--free stars forming in the early minihalos were very massive ($\sim 100~\msun$), owing to the rapid mass accretion enabled by ${\rm H_2}$ cooling. These stars would then leave behind remnant BHs with similar masses \citep{Heger+03}, and produce X-rays, either by direct accretion or by forming high-mass X-ray binaries. A soft X-ray background at photon energies of $\gsim 1$keV, at which the early intergalactic medium (IGM) is optically thin, then provides further global feedback: both by heating the IGM, and by catalyzing ${\rm H_2}$ formation in collapsing halos \citep{HRL96,Oh01,Venkatesan+01,GB03,Madau+04,CM04,Ricotti+05,Mirabel+11}. Recent simulations have been pushed to higher spatial resolution, and in some cases, using sink particles, were able to continue their runs beyond the point at which the first ultra-dense clump developed. The gas in the central regions of at least some of the early minihalos were found to fragment into two or more distinct clumps \citep{Turk+09,Stacy+10,Greif+11,Clark+11,Prieto+11}. This raises the possibility that the first stars formed in multiple systems, and that many of these stars had lower masses than previously thought (but see \citealt{Turk+12} for still higher resolution simulations that suggest less efficient fragmentation). There is also some observational evidence suggesting a lower characteristic Pop III mass. Massive ($\gsim 140~{\rm M_\odot}$), non-rotating metal-free stars are expected to end their lives as pair-instability supernovae (PISNe), and the non--detection of the characteristic PISN nucleosynthetic patterns in metal--poor stars suggests that the typical Pop III stars did not form with such high masses (see, e.g., a recent review by \citealt{FN11}). The observations of carbon--enhanced metal poor stars may further imply a significant number of Pop III stars with masses as low as $M= 1-8~\msun$ \citep{Tumlinson07a,Tumlinson07b}. Finally, the recent discoveries of extremely metal poor stars with no sign of C or N enhancement shows that low-mass star formation could occur at metallicities much lower than previously assumed \citep{Caffau+2011,Caffau+2012}, likely facilitated by dust fragmentation \citep{Schneider+2012}. Motivated by the above, in this {\em Letter}, we examine radiative feedback from an early cosmic IR background, produced by a population of low--mass stars. Although we focus on low--mass PopIII (i.e., metal-free) stars, our conclusions are more generic, and apply at any cosmic epoch when significant numbers of low--mass PopII stars co-exist with massive PopIII stars. Low-mass stars are expected to have soft spectra, even in the metal-free case \citep{TS00,Marigo+01,Schaerer02}, producing significant radiation at $\sim 1$eV, near the photo-detachment threshold (0.76 eV) of the ${\rm H^-}$ ion. ${\rm H^-}$ is a reactant in the dominant formation channel for ${\rm H_2}$, (${\rm H^- + H \rightarrow H_2 + e^-}$) and its destruction can therefore have a dramatic impact on the thermal evolution of metal--free gas. In fact, it is well known that photo-detachment of ${\rm H^-}$ by cosmic microwave background (CMB) photons kept the ${\rm H_2}$ formation rate in the early Universe very low, until the CMB photons redshifted to lower energies at $z \sim 100$ \citep{Hirasawa+69,HP06}. ${\rm H^-}$ photo-detachment can become globally important again once stars begin to form, if they have soft spectra\footnote{To our knowledge, this point was first noted and discussed by \citet{Chuzhoy07}; see \S~\ref{subsec:others} below.}. The aim of this {\em Letter} is to quantify (i) if and when, due to low--mass stars, ${\rm H^-}$ photo-detachment again became the dominant process to limit ${\rm H_2}$ cooling in the earliest protogalaxies, and (ii) to what extent this may have increased or decreased the net global negative radiative feedback in the early Universe. We focus on the importance of this negative feedback for minihalos (as opposed to the more massive halos that cool even in the absence of molecular hydrogen). In order to accomplish this, we perform ``one--zone'' calculations, following the coupled chemical and thermal evolution of the gas in the presence of a cosmological radiation background, including ${\rm H^-}$ photo-detachment by IR photons and ${\rm H_2}$--photodissociation by LW photons. The rest of this paper is organized as follows. In \S~\ref{sec:model} we describe our chemical and thermal modeling. In \S~\ref{sec:results}, we present our results for the relative importance of IR radiative feedback, with various assumptions about the stellar populations; we also compare our results to previous studies. Finally, in \S~\ref{sec:conclusions} we offer our conclusions. Throughout this paper we adopt a standard ${\rm \Lambda CDM}$ cosmological background model: ($\Omega_{\rm DM}, ~\Omega_{\rm b},~\Omega_{\Lambda},~h$)= (0.233, 0.046, 0.721, 0.701) \citep{Komatsu+11}. | \label{sec:conclusions} The main conclusion of this {\it Letter} is that if the mass--fraction of low--mass (few $\msun$) stars exceeded $\sim 90\%$, then their early IR background radiation dominated over the LW background in suppressing ${\rm H_2}$ formation. This low-mass fraction is comparable to those in present-day IMFs, and it is interesting to note that in this limit, star formation in minihalos could be more efficient than if the early stars were massive, owing to the lack of UV feedback and heating inside halos. The early IR background from the low-mass stars would then exert significant net feedback, and regulate the star-formation history in the early Universe once a fraction $f_\ast =$ a few $\times f_{\rm \ast,LW}$ of baryons were converted into stars. The threshold $f_{\rm \ast,LW}$ is the fraction required for strong LW feedback (from massive stars), which is $\sim 0.3\%$ that required for reionization assuming ${\rm n_\gamma}= 10$ ionizing photons per hydrogen atom. Future investigations of radiative feedback in the early Universe, which include low--mass stars, should therefore include ${\rm H^-}$ photo-detachment. Our results also highlight the need for an accurate calculation of the IR photon output of low-mass stars. | 12 | 6 | 1206.1861 |
1206 | 1206.1050_arXiv.txt | We examine the possibility that very massive stars greatly exceeding the commonly adopted stellar mass limit of $150 M_{\odot}$ may be present in young star clusters in the local universe. We identify ten candidate clusters, some of which may host stars with masses up to $600 M_{\odot}$ formed via runaway collisions. We estimate the probabilities of these very massive stars being in eclipsing binaries to be $\gtrsim 30\%$. Although most of these systems cannot be resolved at present, their transits can be detected at distances of 3 Mpc even under the contamination of the background cluster light, due to the large associated luminosities $\sim 10^7 L_{\odot}$ and mean transit depths of $\sim 10^6 L_{\odot}$. Discovery of very massive eclipsing binaries would flag possible progenitors of pair-instability supernovae and intermediate-mass black holes. | Many observations support the statistical argument that the upper limit to initial stellar masses is $\sim 150 M_{\odot}$ for Pop II/I stars \citep{Figer2005, Zinnecker2007}. However, this common notion is challenged by the recent spectroscopic analyses of \citet{Crowther2010}, in which star clusters NGC 3603 and R136 are found to host several stars with initial masses above this limit, including one star R136a1 with a current mass of $\sim 265 M_{\odot}$. Also, candidate pair-instability supernovae, which require progenitors with masses above $200 M_{\odot}$, have been observed in the low redshift universe \citep{Gal-Yam2009}. Therefore, it is worth exploring methods to confirm the existence of a very massive star (VMS), defined here as a star with a stellar mass significantly greater than the stellar mass limit, i.e. $M \gtrsim 200 M_{\odot}$. Unless the VMS is very close by, it is extremely difficult to spatially resolve the VMS from stars in its vicinity. Indeed, the central component of R136 was once thought to be an extremely massive $\gtrsim 10^3 M_{\odot}$ star \citep{Cassinelli1981}, before \citet{Weigelt1985} resolved it as a dense star cluster via speckle interferometry. As for spectroscopic measurements, verification of a single VMS is further complicated by the fact that the effective temperature $T_{eff}$ of Pop I stars above $10^2 M_{\odot}$ depends very weakly on mass, with $\log (T_{eff}/{\rm K})\approx 4.7$--$4.8$ \citep{Bromm2001} for stars between $10^2$--$10^3 M_{\odot}$. Moreover, a hot evolved star with an initial mass below $10^2 M_{\odot}$ can nevertheless reach these temperatures in its post main-sequence evolution and mimic a VMS. The most accurate method of constraining the stellar masses of distant stars is by measuring the radial velocity and light curves of the star in an eclipsing binary \citep{Bonanos2009, Torres2010}. The light curve provides a wealth of information about the binary, including its orbital period, inclination, eccentricity, as well as the fractional radii and flux ratio of the binary members. The radial velocities found from a double-lined spectroscopic binary further provide the mass ratio of the binary. With the above information, the individual masses of each star in the binary can be calculated via Kepler's third law. Searches for massive eclipsing binaries in star clusters within our own Galaxy are already underway \citep{Koumpia2011}, and techniques have been suggested for binary searches in other galaxies \citep{Bonanos2012}. In this {\it Letter}, we estimate the masses and properties of VMSs that may have formed via collision runaways in a number of very young, dense, and massive star clusters in the local universe. We calculate the probability of these VMSs to be in eclipsing binaries, and find their expected transit depths and observability. | A search for periodic flux variations (as shown in Fig. 1) due to transits of the VMS candidates in Table \ref{TableOfClusterAndVMSMass} would be of considerable interest. Although \citet{Crowther2010} made robust arguments against R136a1 being a wide separation binary or an equal-mass binary, this source could still involve a short-period, unequal-mass binary system. The Arches cluster is observed to have no stars currently above the $150 M_{\odot}$ mass limit, but \citet{Crowther2010} also found with contemporary stellar and photometric results that the most luminous stars in the Arches cluster had initial masses approaching $200 M_{\odot}$. The radii of VMSs are dependent on their metallicities and rotation \citep{Langer2007}. If the VMS radii in Table \ref{TableOfClusterAndVMSMass} were smaller by $\sim 25\%$ (e.g. at much lower metallicities), all listed eclipsing probabilities would still remain above 1/3, but the expected transit depth would increase up to $\langle\delta\rangle\sim 20\%$. As for the companion star, for most O stars, the point of unity Thomson optical depth occurs close to the hydrostatic radius, but when stellar mass loss exceeds $\sim 10^{-5} M_{\odot}$ yr$^{-1}$, the photosphere $\tau\sim 1$ occurs in the wind itself, effectively increasing the star's radius. This occurs for Wolf-Rayet companions \citep{Lamontagne1996} and for companions $\gtrsim 60 M_{\odot}$ \citep{Vink2000}, in which case our eclipse probabilities and transit depths are too conservative. Binaries can be broadly classified into detached systems, where neither component fills its Roche lobe, versus semi-detached or over-contact systems, where at least one component exceeds its Roche lobe. VMSs in detached binaries have much more sharply defined eclipses, and more importantly, they do not undergo mass transfer and lose mass to their companion. To find the probability that our VMS candidates in Table \ref{TableOfClusterAndVMSMass} are detached eclipsing binaries, we limit the integration in equation (\ref{FinalEclipseProbability}) to periods $p \gtrsim 5$ days, corresponding to orbital distances where the Roche lobe of the VMS is always greater than its radius \citep{Eggleton1983}. For our VMS candidates, the detached eclipsing binary probability is $\approx 17 \%$, i.e. roughly half of all eclipsing systems. However, non-pristine massive stars can also lose mass via strong winds driven by radiation pressure, with a mass loss rate increasing with metallicity. Post main-sequence VMSs can also lose mass eruptively or via pulsational instabilities, although mass loss near the end of the star's life (e.g. the pulsational pair-instability) is not likely to change the observability of our VMS candidates. Under extraordinary mass loss via winds, \citet{Glebbeek2009} found the highest mass attained by a collision runaway product to be $\sim 400 M_{\odot}$, although the star remained at this mass range for only $\sim 0.2$ Myr. On the contrary, \citet{Suzuki2007} found that stellar mass loss does not inhibit the formation of a VMS of $\sim 10^3 M_{\odot}$. If VMSs do in fact form via collision runaways in young, dense star clusters, and retain sufficient masses at the end of their lives, they may explode as pair-instability supernovae (PISNe) \citep{Yungelson2008}. The creation rate of runaway products is in fact consistent with the current observed PISN rate \citep{Pan2012}. However, the most massive VMSs may collapse directly into an intermediate mass black hole (IMBH) via the photodisintegration instability \citep{Woosley2002}. Tentative evidence has been claimed for IMBHs at the center of old globular clusters \citep{Lou2012}, and extragalactic ultraluminous x-rays sources associated with young star clusters \citep{Ebisuzaki2001, Farrell2009}. The identification of VMSs that can serve as the progenitors of PISNe and IMBHs will help move these extreme astrophysical objects from the realm of speculation into reality. | 12 | 6 | 1206.1050 |
1206 | 1206.1266_arXiv.txt | We present the results from a {\it Chandra} pilot study of 12 massive galaxy mergers selected from Galaxy Zoo. The sample includes major mergers down to a host galaxy mass of 10$^{11}$~$M_\odot$ that already have optical AGN signatures in at least one of the progenitors. We find that the coincidences of optically selected active nuclei with mildly obscured ($N_H \lesssim 1.1 \times 10^{22}$~cm$^{-2}$) X-ray nuclei are relatively common (8/12), but the detections are too faint ($< 40$ counts per nucleus; $f_{2-10~keV} \lesssim 1.2 \times 10^{-13}$~erg~s$^{-1}$~cm$^{-2}$) to reliably separate starburst and nuclear activity as the origin of the X-ray emission. Only one merger is found to have confirmed binary X-ray nuclei, though the X-ray emission from its southern nucleus could be due solely to star formation. Thus, the occurrences of binary AGN in these mergers are rare (0--8\%), unless most merger-induced active nuclei are very heavily obscured or Compton thick. | \label{sec:intro} Major mergers are a key component of current models for galaxy formation in a $\Lambda$CDM Universe. % Mergers can disrupt the star-forming gas and stellar disks of the progenitors, trigger a powerful burst of star formation, and reshape the remaining stellar content into a bulge. Perhaps with a small time delay, the supermassive black holes may feed on gas from the destabilized or destroyed disk, injecting energy in the form of radiation or kinetic outflows that sweep the remnant clear of dust and gas. First proposed by \citet{sanders88}, this picture directly links the triggering of active galactic nuclei (AGNs) phases to the destructive potential induced by a merger. Recent semi-analytic models and hydrodynamic simulations have adopted this scenario to explain the fueling of AGNs and the red spheroidal remnants that are difficult to reproduce without some kind of ``AGN feedback'' \citep{springel, dimatteo, hopkins06, hopkins08, somerville}. In principle, major mergers carry with them two black holes, both of which may be accreting and be visible as distinct AGN during a phase of abundant gas availability that a major, gas-rich merger represents. Yet the evidence associating AGN phases with major mergers remains contested \citep{derobertis, malkan, schmitt, pierce, georgakakis, gabor, schawinski11}. Large optical surveys using Sloan Digital Sky Survey (SDSS) data have found $\sim$3.6\% of spectroscopically confirmed AGNs are in closed binaries \citep[$\sim$5--100~kpc separation;][]{liu11}. The DEEP2 survey also found that binary AGN exist in $\sim$2.2\% (2/91) of red galaxies with type 2 Seyfert optical spectra \citep{deep2a, deep2b}. However, optical surveys can easily miss obscured AGNs especially in merger systems where the gas is driven toward the center through dissipation \citep[e.g.,][]{hopkins08}. X-ray surveys are needed to identify the more highly obscured systems ($N_H \gtrsim 10^{20}$~cm$^{-2}$). We know of only a small number of binary AGN resolved directly using X-ray observations \citep[e.g.,][]{komossa, guainazzi, hudson, bianchi, foreman, comerford, fabbiano}. The intrinsic frequency of binary AGN phases has not been observationally constrained, as the separation of individual X-ray sources is not possible in high-redshift sources and there has been no systematic search for such systems in known mergers. A study of the host galaxies of 185 nearby ($z \lesssim 0.05$) BAT AGNs by \citet{koss11} found that these hard X-ray selected AGNs are preferentially found in massive galaxies with large bulge-to-disk ratios and large supermassive black holes. This may imply that the frequency of binary AGNs is higher in massive mergers. In order to quantify the intrinsic frequency of double AGNs in the local universe, we embarked on a study of the presence of binary AGNs and their dependence on the mass of the host galaxies. The results from the present survey represent a pilot effort as the sample is comprised of only the most massive galaxies in the Galaxy Zoo merger sample. Thus, this paper aims to quantify the intrinsic frequency of double AGNs in the mass limit down to $\sim$10$^{11}$~$M_\odot$ using a study of 12 merging galaxies with the {\it Chandra} X-ray observatory. Throughout this paper, we adopt $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_M = 0.3$, and $\Omega_\Lambda = 0.7$. | \label{sec:disc} The shapes of the X-ray spectra differ for AGNs, obscured AGNs, starbursts, and AGN-star forming composites. Typically, unobscured AGNs have spectra that are well-represented by a power law with photon index of $\sim$1.8. Obscuration affects the lower energy ($\lesssim$2~keV) photons more readily than the higher energy photons and thus flatten or harden the AGN spectra. Starburst spectra are dominated by emission in the lower energies, but low-mass X-ray binaries tend to have relatively flat spectra. Composite objects generally have softened spectra compared to simple AGN spectra due to the soft-energy contribution of the starburst. Given that these mergers contain optically selected AGNs, it is unsurprising that eight of the 10 nuclei have HRs that are consistent with the canonical spectral shape of unobscured AGNs ($\Gamma \sim 1.7 - 2.1$). As many as five could be steeper (GZ~1E, GZ~4S, GZ~5S, GZ~7SW, and GZ~9S), as if star formation is a significant contributor, though the errors in HR and $\Gamma$ allow for unobscured AGN values. Similarly, three nuclei (GZ~2N, GZ~9N, and GZ~11S) have nominally flat spectra, implying dominance from star formation or obscured nuclear activity. However, the errors in HR cannot rule out unobscured AGN as the source of the X-ray emission. Finally, two nuclei (GZ~3S and GZ~10S) have flat or inverted photon indices ($\Gamma \lesssim 1.45$ after accounting for the measurement errors), suggesting some level of obscuration. If we assume a power law with $\Gamma$ fixed at 1.8, the observed HRs imply column densities ($N_H$) $\lesssim 10^{21-22}$~cm$^{-2}$ (Table~\ref{tab:sample}). At these column densities, the HR estimates of the 2--10~keV luminosity are reliable to within $\sim$40\% \citep{teng10}. These columns do not suggest the presence of Compton-thick nuclei though there remains a possibility of leaky, Compton-thick absorbers. \subsection{Starburst Contamination} Of the 12 mergers in the sample, one has no X-ray detection (GZ~8) and only one (GZ~9) exhibits binary X-ray nuclei (Figure~\ref{fig:colorfig}). The remaining 10 mergers have one detected nucleus each. In GZ~5, GZ~7, and GZ~12, the X-ray-detected nucleus is not the one with an optical AGN classification, so in that sense they are double nuclei. In addition, the detected southern nucleus of GZ~10 has extended soft X-ray emission (Figure~\ref{fig:colorfig}), suggesting a contribution from star formation. This raises the question of whether more of the detections might be contaminated by star formation.% To explore this possibility, we compare star formation rates derived from the SDSS {\it u} band luminosities following \citet{sdss} with those derived from the 2--10~keV luminosity following \citet{ranalli} in Figure~\ref{fig:sfrplot}. When compared with the SDSS $u$-band derived star formation rates (Figure~\ref{fig:sfrplot}), the X-ray derived star formation rates of four nuclei (GZ~1E, GZ~5S, GZ~7SW, and GZ~9S) have unconstrained lower limits. While the nominal X-ray derived star formation rates are above the line of equality implying the presence of AGNs, we cannot rule out the possibility that the X-ray emission can be accounted for solely by star formation in these four nuclei. The X-ray luminosities of the remaining nuclei are above those expected from star formation even after the consideration of the $\lesssim$40\% error in the calculation of the X-ray luminosity, consistent with additional contribution to the X-ray luminosity by nuclear activity. Accounting for the error bars, the southern nucleus in GZ~9 may also be dominated by star formation, suggesting GZ~9 does not contain an AGN pair. % \subsection{Compton-thick Nuclei} Three nuclei of the 12 SDSS-selected mergers are not detected in X-rays (GZ~5N, GZ~7NE, and GZ~12E). We already know these have optically identified AGN components, so it is unclear whether the non-detections are due to faint AGNs (two of the three have the highest redshifts in our sample) or Compton-thick AGNs. If we assume these are faint AGNs, a power law model with $\Gamma$ = 1.8 and mild absorption from the Milky Way places upper limits to the luminosity of these objects. Not accounting for intrinsic absorption, the 2-10~keV luminosity for GZ~5N is $\lesssim 5 \times 10^{39}$~erg~s$^{-1}$~cm$^{-2}$ and $\lesssim 5 \times 10^{40}$~erg~s$^{-1}$~cm$^{-2}$ for GZ~7NE and GZ~12E. In the case of the Compton-thick AGNs, the optical signature is coming from the much larger scale narrow- and broad-line regions while the X-ray is sensitive to the small scale emission from the black hole itself. The presence of undetected obscured nuclei would affect our statistics of the frequency of binary AGNs. It is unlikely that all of the secondary nuclei contain Compton-thick X-ray sources, unless an obscured phase is common to mergers (unlike isolated AGNs). Even without a merger-induced obscured phase, the number of heavily obscured AGNs is comparable to the number of less obscured AGNs \citep{tuv}; the presence of Compton-thick nuclei remains a possibility. While the individual detected nuclei have too few counts for spectral fitting to definitively establish whether Compton-thick AGNs are present, we considered the cumulative rest-frame photon distribution of the detected nuclei in the hard band. We compared this observed distribution with the expected photon distributions from unobscured AGNs and from Compton-thick AGNs. In the former case, we assumed a single unabsorbed power law; in the latter case, we assumed a power law with an iron emission line at 6.4~keV with an equivalent width of 1~keV, a typical signature of Compton-thick AGNs. For both cases, the total photon counts were normalized to be the same as the total detected counts. In Figure~\ref{fig:cumuplot} we plot the cumulative distribution of the detected photons in our sample. There is no clear distinction between the observed distribution with either model. In fact, the two-tailed Kolmogorov-Smirnov (K-S) test statistics for the two cases are nearly identical. As a sanity check, we compared the two modeled distributions with each other and there is a clear difference at the $\sim$80\% confidence level. Therefore, we cannot rule out the possibility that Compton-thick AGNs are present at the level that we are able to detected these sources. \subsection{U/LIRGs in Formation?} In theoretical models of galaxy mergers \citep[e.g.,][]{hopkins08}, luminous and ultraluminous infrared galaxies (U/LIRGs) represent a stage that mergers go through before the formation of elliptical galaxies. Initially, tidal torques enhance star formation and black hole accretion. Then in the final coalescence of the galaxies, massive inflows of gas trigger starbursts with strengths similar to those inferred for U/LIRGs. The mergers in our sample appear to be the predecessors to U/LIRGs in this evolutionary picture. The X-ray luminosities estimated for our mergers are approximately 10 times lower than those observed in most U/LIRGs, but are consistent with the lower end of the range measured in LIRGs \citep{teng10, lehmer, goals}. This implies mergers in our sample are in the earliest stages of interaction, where the growth of the central black hole has not yet peaked. The incidence of binary AGNs in U/LIRGs is also rare. The Revised Bright Galaxy Survey \citep[RBGS;][]{rbgs} is a flux-limited sample of U/LIRGs from the {\it IRAS} All Sky Survey. Of the 629 extragalactic objects with 60~$\mu$m flux greater than 5.24~Jy, 86 are interacting galaxies that are visually similar to our sample in the optical (i.e. close binaries). Of these, 32 have high-quality X-ray data from either {\it Chandra} or {\it XMM-Newton} that is sensitive to the presence of an AGN. Not accounting for the presence of undetected Compton-thick nuclei, only 3\% (1/32) of the RBGS sources with X-ray data show binary X-ray nuclei \citep[NGC~6240;][]{komossa}. This is consistent with the 0--8\% (0--1 out of 12) we observe in our modest SDSS sample. \subsection{Frequency of Binary AGNs in SDSS Mergers} From the very short snapshots of our study, we have found that coincidence of optically selected active nucleus with mildly obscured X-ray nucleus is relatively common (8/12). Given the faint detections, these snapshots are too short to place strong limits on the absence of AGN in the undetected galaxies, so it is difficult to comment on the frequency of binary active nuclei. However, we do detect a pair of X-ray nuclei in GZ~9, implying that this is uncommon unless the second nucleus is heavily obscured. In that instance, the most likely scenario would be that all nuclei are obscured. That is, either binary nuclei are uncommon, or merger nuclei in general have a high probability of being heavily obscured. The latter possibility cannot be addressed by the current sample. To do better, we will need to increase the exposure times, expand our merger sample for better statistics, and include a sample of major mergers for which there are no optically detected nuclei. Another natural follow-up would be to extend the study to a similarly selected sample with a lower mass limit to examine the dependence of binary AGNs on the mass of the host galaxies. | 12 | 6 | 1206.1266 |
1206 | 1206.5786_arXiv.txt | Inspiraling supermassive black hole binary systems with high orbital eccentricity are important sources for space-based gravitational wave (GW) observatories like the Laser Interferometer Space Antenna (LISA). Eccentricity adds orbital harmonics to the Fourier transform of the GW signal, and relativistic pericenter precession leads to a three-way splitting of each harmonic peak. We study the parameter estimation accuracy for such waveforms with different initial eccentricity using the Fisher matrix method and a Monte Carlo sampling of the initial binary orientation. The eccentricity improves the parameter estimation by breaking degeneracies between different parameters. In particular, we find that the source localization precision improves significantly for higher-mass binaries due to eccentricity. The typical sky position errors are $\sim1\,$deg for a nonspinning, $10^7\,M_{\odot }$ equal-mass binary at redshift $z=1$, if the initial eccentricity 1 yr before merger is $e_0\sim 0.6$. Pericenter precession does not affect the source localization accuracy significantly, but it does further improve the mass and eccentricity estimation accuracy systematically by a factor of 3--10 for masses between $10^6\,M_{\odot }$ and $10^7\,M_{\odot }$ for $e_0 \sim 0.3$. | The inspiral and merger of compact binary systems of black holes are important sources of gravitational waves (GWs) for the proposed space-based GW missions such as the Laser Interferometer Space Antenna (LISA) \cite{LISA1} or the European New Gravitational Wave Observatory (NGO/eLISA) \cite{eLISA1}. The detectable frequency band for these instruments will be around $10^{-4}$ to $10^{-1}\,$Hz \cite{LISA2} which corresponds to the inspiral of two $(10^{4}-10^{7})M_{\odot }$ black holes. As the sources detected by LISA/NGO will be loud with a large signal-to-noise ratio in general, an ideal method for parameter extraction is matched filtering \cite{Matched}. An effective matched filtering requires an accurate model of the emitted GWs. In this technique the detected signal output is cross-correlated with theoretical waveform templates. In particular, matched filtering is sensitive to the phase information of the waveform, and a high correlation between the signal and template allows one to make predictions on the source parameters \cite{FC,CF}. Many previous studies in the literature adopted waveforms generated by binaries in circular orbits (see Refs. \cite{Thorne,Cutler98,Vecchio,BBW,LH,Arun,TS,LH08,McWilliams, HH,Bence-apj06,Bence-prd,Bence-apj08} for LISA parameter estimation). This is due to the expectation that the orbit of the binary will circularize due to the emission of GWs \cite{PM,P}. Nevertheless, there are a number of reasons to expect that at least some LISA sources may be eccentric. If the binary is embedded in a gaseous disk, it can remain eccentric until the final year of the inspiral \cite{AN,MacFadyen,Cuadra,Sesana}. The interaction of the supermassive black hole (SMBH) binary with a population of stars also increases its eccentricity \cite{Preto,Matsubayashi,Lockmann}. The eccentricity can also be excited by the Kozai mechanism and relativistic orbital resonances in hierarchial triples \cite{Wen,HL,Seto2,AP,Naoz} or by a triaxial potential \cite{HattoriYoshii,MerrittVasiliev}, and may be typical for extreme mass ratio inspirals \cite{AmaroSeoane07,AmaroSeoane12}. Further, black hole binaries in dense galactic nuclei formed by GW emission during close encounters remain very eccentric until merger \cite{Bence,KL}. Population synthesis and binary evolutionary models show that a fraction of stellar compact object binaries may also be eccentric for ground-based (Advanced LIGO/VIRGO and Einstein Telescope) and space-based detectors (DECIGO) \cite{Kowalska}. Including eccentricity in the waveform may be essential for the detection of inspiraling eccentric binaries with matched filtering and to avoid a systematic bias in the parameter estimation \cite{CV}. Using circular templates to detect waveforms with eccentricities $e_0\gtrsim 0.1$, leads to a significant loss of signal-to-noise ratio for ground-based detectors such as LIGO and VIRGO \cite{MP,BrownZimmer}. A similar conclusion was reached for eccentric massive black hole binaries detected with LISA \cite{PorSe}. The orbital evolution and waveforms have been developed to first and second post-Newtonian (PN) order, including spin-orbit and spin-spin contributions for eccentric orbits~\cite{KMG,KleinJetzer,GopakumarSchaefer,TessmerSchaefer,Hinder,CornishKey}. To assess the astrophysical impact of planned GW instruments, it is essential to estimate the expected parameter measurement precision of typical GW sources. This may be done by injecting a simulated GW signal into synthetic detector noise and carrying out a Monte Carlo Markov Chain (MCMC)-based matched filtering search for a parametrized template model to recover the posterior distribution function (PDF) of the estimated source parameters \cite{CP}. Porter and Sesana \cite{PorSe} investigated the cases of low-mass ($100M_{\odot }$) and high ($10^{4}M_{\odot }$) mass black hole binaries on eccentric orbits using nonspinning, restricted 2 PN waveforms. They concluded that eccentricity can significantly bias the recovered parameters of the source for LISA if circular templates are used even if the eccentricity is as small as $e \sim 10^{-4}$. More recently, Key and Cornish \cite{KeyCornish} extended that study by using an effective 1.5PN waveform for inspiraling eccentric SMBHs [with $m\sim(10^{5}-10^{7})M_{\odot}$] taking into account eccentricity and spin effects in the template model. They found that the eccentricity measurement errors are of order $\Delta e \sim 10^{-3}$ for a range of mass ratios and a particular choice of angular parameters. Since the parameter space is large, $17$-dimensional for an eccentric spinning binary, state-of-the-art MCMC calculations are numerically too expensive to explore the full range of source parameters. However, for a large signal-to-noise ratio (SNR), the PDF may be well approximated by an ellipsoid, and the parameter measurement errors can be estimated very efficiently using the Fisher matrix method \cite{FC,CV}. Using this method, it has been shown that different source inclinations and sky locations lead to a wide range of parameter measurement errors subtending many orders of magnitude \cite{Hughes,Vecchio,BBW,LH,LH08}. In this study, we carry out a Fisher matrix analysis to investigate the possible range of parameter estimation errors for eccentric binaries. Only a few studies have investigated the LISA parameter estimation accuracy for eccentric inspiraling sources using the Fisher matrix method (cf. Refs. \cite{Cutler98,Hughes,Vecchio,BBW,LH,LH08,Arun,TS} for circular inspirals). Barack and Cutler \cite{BC} investigated the LISA errors for highly eccentric stellar mass compact objects inspiraling into a SMBH. They found that the influence of eccentricities on $\Delta \mathcal{M}/\mathcal{M}\sim 10^{-4}$ (error of the chirp mass), $\Delta e_{0}\sim 10^{-4}$ (error of initial eccentricity) and $\Delta \Omega_{S}\sim 10^{-4}$ (angular resolution error) is not substantial; the error estimates do not differ much from those obtained for circular orbits \cite{Cutler98}. However, they assumed only an arbitrarily chosen, single set of orientations, which may not be representative of the typical errors. Yunes et al. \cite{Yunes} provided ready-to-use analytic expressions for the Fourier waveforms of moderately eccentric sources. They have shown that eccentricity increases the detectable mass range of GW detectors toward higher masses by enhancing the orbital harmonics \cite{Arun,TS}. Yagi and Tanaka \cite{Tanaka} investigated the LISA errors for various alternative theories of gravity for spinning, small-eccentricity inspiraling SMBH binaries ($e_0\sim 0.01$ at 1 yr before merger), using restricted 2 PN waveforms, neglecting higher orbital harmonics and apsidal precession in the waveform. They have found that the eccentricity and the spin-orbit interaction reduce the parameter errors by an order of magnitude for spinning SMBHs in massive graviton theories, but not in Brans-Dicke-type theories. Neither of the previous systematic Fisher matrix studies of parameter errors included the effects of relativistic pericenter precession for eccentric sources. However, precession effects introduce an additional feature in the waveform, and have the potential to break the degeneracy between parameter errors \cite{Bence-prd}. In particular, spin-orbit precession has been shown to improve the source localization precision substantially during the last day of the inspiral \cite{Vecchio,LH,LH08}. Similarly, GR pericenter precession may also be expected to improve the LISA parameter measurement accuracy. In fact, since pericenter precession enters at a lower PN order, this improvement could take place well before the binary reaches merger. Localizing the source before merger could be used to provide triggers for electromagnetic (EM) facilities to search for the EM counterpart \cite% {Bence-apj08}. A coincident GW and EM observation of the same source could have far-reaching astrophysical implications \cite{Schutz,HH,Bence-apj06,Bence-apj08} In the present paper, we carry out a systematic parameter estimation study for inspiraling SMBH binaries, taking into account both orbital eccentricity and the relativistic pericenter precession effect. We account for the evolution of the semimajor axis and eccentricity in our waveforms to leading order due to GW emission \cite{Whalquist,MP,Pierro,BC}, but we neglect higher-order PN contributions and spin effects. We compute the waveform in the frequency domain using the stationary phase approximation (SPA, see Refs. \cite{Yunes,Tessmer,ecc1,ecc2,Seto}) and derive the signal-to-noise ratio (SNR) and the Fisher information matrix using a Fourier-Bessel analysis for the parameter estimation of eccentric sources. To explore the possible range of parameter errors, we generate a Monte Carlo sample of binaries with random orientations and vary the masses and initial eccentricities systematically over a wide range relevant for LISA. We calculate the parameter errors for the standard three-arm LISA/NGO configuration, as well as for a descoped detector configuration, where one of the two independent interferometers is removed. In Sec.~II, we summarize the basic formulas describing eccentric waveforms in the leading quadrupole approximation, using a Fourier-Bessel decomposition. In Sec.~III, we derive the frequency domain waveforms and the LISA detector response. After a brief introduction of parameter estimation using the Fisher matrix method in Sec. IV, we present results for specific systems in Sec. V. We summarize our conclusions in Sec VI. Some details of the calculations are described in Appendixes A and B. We use geometrical units $G=c=1$. | We carried out an extensive study of parameter estimation for eccentric binaries with arbitrary orbital eccentricity. We computed the waveforms in the frequency domain by a new method optimized for taking into account eccentricity, by changing the integration variable for the waveforms from the orbital frequency $\nu(e)$ to the eccentricity variable $e$ \cite{Bence}. This results in an improvement of numerical precision as compared to standard approaches in the frequency domain, where a Taylor series expansion of the orbital frequency $\nu(e)$ (among others) in the eccentricity $e$ is needed \cite{Yunes}. Our method is well suited for computing the Fisher matrix and the signal-to-noise ratio. Our parameter space is ten dimensional, consisting of four angles, the chirp mass, the luminosity distance, coalescence time and phase, initial eccentricity and pericenter position (compare Fig. \ref{geometry}). The first eight parameters are standard for circular orbits too. We have examined the LISA parameter estimation errors for GWs emitted by eccentric inspiraling SMBH binaries including the effects of pericenter precession. Based on a large set of simulated binary waveforms, we found that there is about 1 order of magnitude improvement compared to circular waveforms in LISA's angular resolution for highly eccentric sources (e.g. $e_{0}=0.6$) for relatively high SMBH masses $\sim 10^{7}M_{\odot }$. There is however, a much smaller effect for lower-mass binaries in the range $(10^{4}-10^{5})M_{\odot }$. This improves the prospects for identifying the electromagnetic counterparts \cite{Bence-apj06,Bence-apj08} of relatively high-mass eccentric SMBH mergers with LISA. Similar conclusions have been reached in Refs.~\cite{Arun,TS}. However, we found that pericenter precession does not further improve the sky localization accuracy of the source, although it may further improve the measurement errors of mass and eccentricity parameters. It is important to note that the angular resolution is significantly affected by the number of detectors (see Figs. \ref{angular66} and \ref{angular77}). However, nearly the same parameter estimation accuracy can be obtained for the single and total detector configurations for $(10^6-10^6)M_{\odot }$ binaries for fast parameters \citep{Bence-prd} like the chirp mass and eccentricity (Figs. \ref{chirp66} and \ref{e066}). The second detector systematically reduces the errors of these parameters for higher masses $(10^7-10^7)M_{\odot }$. \begin{table*}[tb] \caption{Parameter estimation errors for equal-mass SMBH binaries. The initial eccentricities $e_{0}$ are $10^{-6}$ (nearly circular), $0.3$, and $0.6.$, the luminosity distance is $D_{L}=6.4\mathrm{Gpc}$ ($z=1$); and the angular parameters are $\protect\phi _{L}=4.724$, $\protect\mu _{L}=-0.3455$ , $\protect\phi _{S}=4.642$, and $\protect\mu _{S}=-0.3185$.} \label{table2} \begin{center} \begin{tabular}{c|c|ccccc} \hline\hline $\underset{(M_{\odot })}{SMBH}$ & $e_{0}$/precession & $SNR$ & $\underset{ (\times 10^{-2})}{\Delta D_{L}/D_{L}}$ & $\underset{(\times 10^{-6})}{\Delta \mathcal{M}/\mathcal{M}}$ & $\underset{(\times 10^{-6})}{\Delta e_{0}}$ & $\underset{(\times 10^{-6})}{\Delta \Omega }$ \\ \hline\hline $10^{7}-10^{7}$ & \begin{tabular}{l} \begin{tabular}{l} $e_{0}=10^{-6}$, no prec. \\ $e_{0}=10^{-6}$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.3$, no prec. \\ $e_{0}=0.3$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.6$, no prec. \\ $e_{0}=0.6$, incl. prec. \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1119$ \\ $2002$% \end{tabular} \\ \begin{tabular}{c} $1116$ \\ $1984$% \end{tabular} \\ \begin{tabular}{c} $1146$ \\ $1984$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $837$ \\ $538$ \end{tabular} \\ \begin{tabular}{c} $96.2$ \\ $42.9$% \end{tabular} \\ \begin{tabular}{c} $31.6$ \\ $17.3$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $105$ \\ $9.14$% \end{tabular} \\ \begin{tabular}{c} $67.7$ \\ $9.42$ \end{tabular} \\ \begin{tabular}{c} $17.4$ \\ $4.95$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1794$ \\ $1311$ \end{tabular} \\ \begin{tabular}{c} $222$ \\ $34.7$ \end{tabular} \\ \begin{tabular}{c} $6.91$ \\ $2.14$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $193$ \\ $77.9$ \end{tabular} \\ \begin{tabular}{c} $3.32$ \\ $0.893$% \end{tabular} \\ \begin{tabular}{c} $2.16$ \\ $0.689$ \end{tabular} \end{tabular} \\ \hline $10^{6}-10^{6}$ & \begin{tabular}{l} \begin{tabular}{l} $e_{0}=10^{-6}$, no prec. \\ $e_{0}=10^{-6}$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.3$, no prec. \\ $e_{0}=0.3$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.6$, no prec. \\ $e_{0}=0.6$, incl. prec. \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1171$ \\ $1704$ \end{tabular} \\ \begin{tabular}{c} $1176$ \\ $1701$ \end{tabular} \\ \begin{tabular}{c} $1200$ \\ $1712$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $192$ \\ $168$ \end{tabular} \\ \begin{tabular}{c} $30.6$ \\ $26.0$% \end{tabular} \\ \begin{tabular}{c} $10.3$ \\ $8.29$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $3.09$ \\ $1.19$ \end{tabular} \\ \begin{tabular}{c} $3.99$ \\ $1.51$ \end{tabular} \\ \begin{tabular}{c} $3.17$ \\ $1.56$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1562$ \\ $1363$ \end{tabular} \\ \begin{tabular}{c} $7.53$ \\ $3.32$ \end{tabular} \\ \begin{tabular}{c} $1.18$ \\ $0.917$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $13.5$ \\ $9.33$ \end{tabular} \\ \begin{tabular}{c} $2.00$ \\ $1.00$ \end{tabular} \\ \begin{tabular}{c} $1.84$ \\ $0.871$ \end{tabular} \end{tabular} \\ \hline $10^{5}-10^{5}$ & \begin{tabular}{l} \begin{tabular}{l} $e_{0}=10^{-6}$, no prec. \\ $e_{0}=10^{-6}$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.3$, no prec. \\ $e_{0}=0.3$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.6$, no prec. \\ $e_{0}=0.6$, incl. prec. \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1924$ \\ $2183$ \end{tabular} \\ \begin{tabular}{c} $1925$ \\ $2184$ \end{tabular} \\ \begin{tabular}{c} $1920$ \\ $2188$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $314$ \\ $296$ \end{tabular} \\ \begin{tabular}{c} $33.4$ \\ $26.6$ \end{tabular} \\ \begin{tabular}{c} $14.3$ \\ $12.0$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $1.03$ \\ $0.958$ \end{tabular} \\ \begin{tabular}{c} $1.30$ \\ $1.16$ \end{tabular} \\ \begin{tabular}{c} $1.04$ \\ $1.23$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $2595$ \\ $2365$ \end{tabular} \\ \begin{tabular}{c} $2.74$ \\ $3.54$ \end{tabular} \\ \begin{tabular}{c} $0.435$ \\ $0.831$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $30.6$ \\ $25.6$ \end{tabular} \\ \begin{tabular}{c} $0.848$ \\ $0.553$ \end{tabular} \\ \begin{tabular}{c} $0.678$ \\ $0.520$ \end{tabular} \end{tabular} \\ \hline $10^{4}-10^{4}$ & \begin{tabular}{l} \begin{tabular}{l} $e_{0}=10^{-6}$, no prec. \\ $e_{0}=10^{-6}$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.3$, no prec. \\ $e_{0}=0.3$, incl. prec. \end{tabular} \\ \begin{tabular}{l} $e_{0}=0.6$, no prec. \\ $e_{0}=0.6$, incl. prec. \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $306$ \\ $314$ \end{tabular} \\ \begin{tabular}{c} $310$ \\ $318$ \end{tabular} \\ \begin{tabular}{c} $333$ \\ $341$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $746$ \\ $697$ \end{tabular} \\ \begin{tabular}{c} $71.2$ \\ $62.9$ \end{tabular} \\ \begin{tabular}{c} $30.0$ \\ $27.3$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $0.628$ \\ $1.93$ \end{tabular} \\ \begin{tabular}{c} $0.847$ \\ $1.80$ \end{tabular} \\ \begin{tabular}{c} $0.539$ \\ $0.925$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $4605$ \\ $4433$ \end{tabular} \\ \begin{tabular}{c} $1.60$ \\ $4.68$ \end{tabular} \\ \begin{tabular}{c} $0.193$ \\ $0.356$ \end{tabular} \end{tabular} & \begin{tabular}{c} \begin{tabular}{c} $239$ \\ $189$ \end{tabular} \\ \begin{tabular}{c} $30.3$ \\ $29.0$ \end{tabular} \\ \begin{tabular}{c} $28.3$ \\ $27.4$ \end{tabular} \end{tabular} \\ \hline\hline & & & & & & \end{tabular} \end{center} \end{table*} | 12 | 6 | 1206.5786 |
1206 | 1206.2359_arXiv.txt | One of the strongest predictions of the $\Lambda$CDM cosmological model is the presence of dark satellites orbiting all types of galaxies. We focus here on the dynamical effects of such satellites on disky dwarf galaxies, and demonstrate that these encounters can be dramatic. Although mergers with $M_{\rm sat} > M_{d}$ are not very common, because of the lower baryonic content they occur much more frequently on the dwarf scale than for $L_*$-galaxies. As an example, we present a numerical simulation of a 20\% (virial) mass ratio merger between a dark satellite and a disky dwarf (akin to the Fornax dwarf galaxy in luminosity) that shows that the merger remnant has a spheroidal morphology. We conclude that perturbations by dark satellites provide a plausible path for the formation of dSph systems and also could trigger starbursts in gas rich dwarf galaxies. Therefore the transition from disky to the often amorphous, irregular, or spheroidal morphologies of dwarfs could be a natural consequence of the dynamical heating of hitherto unobservable dark satellites. | \label{sec:intro} According to the $\Lambda$CDM scenario, stellar disks are immersed in dark matter halos and are surrounded by a full spectrum of satellite companions. Encounters with these satellites can inject significant amounts of energy into the system, with consequences that vary from negligible to fully catastrophic disk destruction depending on the relative mass of the perturber and the configuration of the event (relative distances and velocities). Disk heating by such substructures has been addressed in previous work \citep{Toth1992,Quinn1993,Font2001,Benson2004}, but has generally focused on the effect on bright Milky Way-like galaxies. Cold dark matter models predict the structure of halos to be self-similar; in such a way that, when properly scaled, a Milky Way-sized halo looks comparable to one hosting a faint dwarf galaxy \citep{Moore1999b,AqPaper2,Klimentowski2010,Wang2012}. However, galaxy formation is not a self-similar process, as the properties of galaxies depend in a complex way on e.g.\ the mass of their host halos. For example, low mass (dwarf) galaxies are much more inefficient at forming stars \citep{Blanton2001,RK2008} and have much higher mass-to-light ratios than larger galaxies \citep{Yang2003,walker2009}. In addition, gas cooling is likely to be (nearly) completely inhibited in dark matter halos with masses below $\sim 10^8 h^{-1} \sm$ \citep{Kaufmann2007}, which implies that the satellites of dwarfs should be generally completely dark in contrast to satellites in galaxy clusters or around $L_*$-galaxies. In this {\it Letter} we show that these considerations imply that the dynamical perturbations of dark-matter satellites on dwarf galaxies are much more important than on $L_*$-galaxies. Dark satellites may provide a channel for the formation of dwarf spheroidal galaxies without the need to recur to environmental effects \citep{Mayer2010} or multiple body interactions \citep{Sales2007}. Such interactions may also be responsible for the observed increase of disk ``thickness'' towards fainter galaxies \citep{Yoachim2006}, as well as explain the existence of isolated dwarfs undergoing intense starbursts without an apparent trigger \citep{Bergvall2011} as a result of a major merger with a dark companion (T.~K.~Starkenburg et al., in prep.). | We have demonstrated that the dynamical effects of dark satellites on disky dwarf galaxies are much more dramatic than on galaxies like the Milky Way. Mergers with $M_{\rm sat} > M_{d}$ are not very common for $z < 3$ but they occur much more frequently than on the $L_*$-galaxies scale. As an example, we have simulated a merger with $M_{\rm sat}/M_{\rm vir} = 0.2$ for a dwarf with $M_d = 8 \times 10^{7} h^{-1} \sm$ in stars, i.e. slightly more massive than the Fornax dwarf galaxy, and found that its morphology changed from disky to spheroidal. This might be a plausible path for the formation of dSph systems in isolation (if the dwarf was gas poor, which is rare in our models but not unlikely). This channel might also be relevant for the dSph satellites of our Galaxy, provided such encounters would have taken place just before the system fell onto the potential well of the Milky Way (since further gas accretion would thus be prevented). Most of the galaxies on the scales of dwarfs are, however, gas-rich. In that case, encounters with dark satellites can trigger starbursts, which might explain the presence of seemingly isolated dwarfs undergoing major star formation events without an apparent trigger. Depending on the characteristics of the encounter, such starbursts will vary in amplitude. We are currently performing hydrodynamical simulations to characterize this process (T.~K.~Starkenburg et al. in prep). Additionally, other processes exist that can influence the morphologies of dwarf galaxies. For example, binary mergers between disky dwarfs can result in the formation of spheroidal systems \citep{stelios2011}, although such events are rare (see Fig.~1). On the other hand, on the scale of dwarf galaxies physical processes affecting gas may also lead to thicker systems. For example, the presence of a temperature floor in the interstellar medium at $T\sim 10^4K$ introduced by e.g. a UV background, implies that gas pressure becomes comparable to rotational support for small dark matter halos. Stars formed in such systems would thus be born in puffier configurations as demonstrated by \citet{Kaufmann2007} \citep[see also, e.g.][]{RK2008}. Yet, we have shown here that a distinctive imprint on dwarf galaxies will be left by dark satellites in the context of the $\Lambda$CDM cosmological paradigm. Such {\it dark satellites} are expected to make the stellar disks of isolated dwarf galaxies significantly thicker than those of $\sim L_*$ galaxies. We have indeed detected such a trend on three different observational samples of isolated late-type galaxies on the nearby Universe. We may have identified a new mechanism to explain the morphologies of dwarf galaxies. | 12 | 6 | 1206.2359 |
1206 | 1206.7056_arXiv.txt | \noindent Recently reported tentative evidence for a gamma-ray line in the Fermi-LAT data is of great potential interest for identifying the nature of dark matter. We compare the implications for decaying and annihilating dark matter taking the constraints from continuum gamma-rays, antiproton flux and morphology of the excess into account. We find that higgsino and wino dark matter are excluded, also for nonthermal production. Generically, the continuum gamma-ray flux severely constrains annihilating dark matter. Consistency of decaying dark matter with the spatial distribution of the Fermi-LAT excess would require an enhancement of the dark matter density near the Galactic center. | Monochromatic gamma-ray lines have been suggested long ago as signature for pair-annihilation of dark matter (DM) particles \cite{Bergstrom:1988fp}. During the past decades many DM candidates have been discussed, the most popular ones being weakly interacting massive particles (WIMPs) in the context of supersymmetric extensions of the Standard Model \cite{Bertone:2004pz}. More recently, also decaying dark matter has been studied in detail \cite{Takayama:2000uz,Buchmuller:2007ui,Cirelli:2012ut,Arvanitaki:2008hq}. An attractive feature of decaying gravitino dark matter is the consistency with thermal leptogenesis, contrary to standard WIMP dark matter \cite{Buchmuller:2007ui}. During the past years the Large Area Telescope (LAT)~\cite{Atwood:2009ez}, on board the Fermi gamma-ray space telescope, has searched with unprecedented sensitivity for photon lines from 30~MeV to 300~GeV. Stringent constraints on decaying and annihilating dark matter have been obtained from searches in the energy range $30-200$~GeV based on 11 months of data~\cite{Abdo:2010nc}, and in the range $7-200$~GeV based on 23 months of data~\cite{Ackermann:2012qk}. The search region used in these analyses covers the whole sky except for the Galactic disk ($|b|>10^\circ$) plus a $20^\circ\times 20^\circ$ region around the Galactic center. A similar analysis has been performed independently based on publicly available data corresponding to 27 months in Ref.~\cite{Vertongen:2011mu}. No indications for gamma-ray lines were found. In a recent analysis, that is based on optimized search regions around the Galactic center, and takes 43 months of Fermi-LAT data into account, a hint for a gamma-ray feature in the energy range $120-140$~GeV is reported~\cite {Bringmann:2012vr,Weniger:2012tx}. When interpreted in terms of a gamma-ray line~\cite{Weniger:2012tx}, the significance obtained from the statistical uncertainties in the search regions close to the Galactic center is $4.6\sigma$. The significance is reduced to $3.3\sigma$ when correcting for the bias introduced by selecting the search regions. While the excess is currently under an active debate~\cite{Profumo:2012tr,Tempel:2012ey,Boyarsky:2012ca, Bergstrom:2012fi,Ibarra:2012dw,Dudas:2012pb}, the claim has been further strengthened by a recent analysis~\cite{Su:2012ft}, which confirms the existence of an excess, as well as its spectral shape, with even higher statistical significance than the one claimed in Ref.~\cite{Weniger:2012tx}. In addition, indications are found that the excess originates from a relatively narrow region of a few degrees around the Galactic center, possibly with a small offset within the Galactic plane. Hopefully, the question about the existence of the spectral feature and its precise properties will be settled in the near future. Assuming that the feature is real, the question about its origin is of great interest. While an astrophysical explanation might eventually be identified, a spectral feature in this energy range can arise rather generically from the annihilation or decay of dark matter particles. In this paper we compare decaying and annihilating dark matter based on two prototype models: decaying gravitinos and wino/higgsino-like WIMPs which annihilate predominantly into two W-bosons. In both cases we treat the branching ratios into $\gamma\nu$ (gravitino) and $\gamma\gamma$ (wino/higgsino) final states as free parameters to account for some model dependence. Note that we do not demand thermal freeze-out for WIMPs. Higgsinos and winos can be nonthermally produced in gravitino decays, compatible with leptogenesis \cite{Ibe:2011aa,Buchmuller:2012bt}, or, alternatively, in moduli decays~\cite{Moroi:1999zb,Kitano:2008tk,Acharya:2012tw}. In the following we shall compare interpretations of the tentative 130~GeV photon line in terms of decaying and annihilating dark matter in a sequence of increasing assumptions: We first consider constraints from continuum gamma-rays, which are independent of charged cosmic rays and the dark matter distribution (Section~2). This is followed by a discussion of constraints from antiprotons, which depend on the propagation model (Section~3). We then analyze the implications of the spatial distribution of the Fermi-LAT excess (Section~4) and draw our conclusions (Section~5). | In light of tentative evidence for a gamma-ray line observed by Fermi-LAT near the Galactic center, we have investigated the consequences for several prototype scenarios of decaying and annihilating dark matter, motivated by supersymmetric models with gravitino, higgsino or wino-like LSP. We find that, independently of the actual dark matter distribution, the consistency of continuum and monochromatic contributions to the photon spectrum from dark matter decay or annihilation requires a branching ratio into monochromatic photons larger than $\BR_\gamma \gtrsim 0.5\%$. Both higgsino and wino dark matter can be ruled out because of a too large continuum flux, independently of the production mechanism, while gravitino dark matter with wino NLSP is compatible. We have also investigated constraints arising from the primary antiproton flux, which depend on the adopted propagation model. The resulting limits are comparable to the ones from continuum gamma-rays for annihilating dark matter, but can be more important for decaying dark matter. The morphology of the tentative excess, if confirmed, could yield valuable information on the distribution of dark matter close to the Galactic center. When assuming conventional NFW or Einasto profile functions, the required partial lifetimes for decaying dark matter $\tau_{\gamma\nu} \sim (1-3)\times 10^{28}$\,s are in tension with lower limits obtained from the Galactic halo. This tension is less pronounced for more cuspy contracted profiles. If decaying gravitino dark matter is indeed responsible for the excess, the dark matter density should be enhanced in the Galactic center region compared to conventional models. | 12 | 6 | 1206.7056 |
1206 | 1206.5053_arXiv.txt | \input{abstract.tex} | \input{introduction.tex} | \input{conclusion.tex} | 12 | 6 | 1206.5053 |
1206 | 1206.5579_arXiv.txt | The exoplanets known as hot Jupiters---Jupiter-sized planets with periods less than 10 days---likely are relics of dynamical processes that shape all planetary system architectures. Socrates et al. (2012) argued that high eccentricity migration (HEM) mechanisms proposed for situating these close-in planets should produce an observable population of highly eccentric proto-hot Jupiters that have not yet tidally circularized. HEM should also create failed-hot Jupiters, with periapses just beyond the influence of fast circularization. Using the technique we previously presented for measuring eccentricities from photometry (the ``photoeccentric effect''), we are distilling a collection of eccentric proto- and failed-hot Jupiters from the \kep Objects of Interest (KOI). Here we present the first, KOI-1474.01, which has a long orbital period (69.7340 days) and a large eccentricity $e = 0.81^{+0.10}_{-0.07}$, skirting the proto-hot Jupiter boundary. Combining \kep photometry, ground-based spectroscopy, and stellar evolution models, we characterize host KOI-1474 as a rapidly-rotating F-star. Statistical arguments reveal that the transiting candidate has a low false-positive probability of 3.1\%. KOI-1474.01 also exhibits transit timing variations of order an hour. We explore characteristics of the third-body perturber, which is possibly the ``smoking-gun'' cause of KOI-1474.01's large eccentricity. Using the host-star's rotation period, radius, and projected rotational velocity, we find KOI-1474.01's orbit is marginally consistent with aligned with the stellar spin axis, although a reanalysis is warranted with future additional data. Finally, we discuss how the number and existence of proto-hot Jupiters will not only demonstrate that hot Jupiters migrate via HEM, but also shed light on the typical timescale for the mechanism. | The start of the exoplanet era brought with it the discovery of an exotic new class of planets: Jupiter-sized bodies with short-period orbits ($P \lesssim 10$~days), commonly known as hot Jupiters \citep{1995M,1997M}. Most theories require formation of Jupiter-sized planets at or beyond the so-called ``snow line,'' located at roughly a few AU,\footnote{\citet{2008KB} and \citet{2008KK} explore in detail the location of the ice line for different stellar and disk parameters.} and debate the mechanisms through which they ``migrated'' inward to achieve such small semimajor axes. The leading theories fall into two categories: smooth migration through the proto planetary disk \citep[e.g.][]{1980G,1997W,2005AM,2008I,2011BK}, or what \citet{2012SK} (hereafter S12) term high eccentricity migration (HEM), in which the planet is perturbed by another body onto an inclined and eccentric orbit that subsequently circularizes through tidal dissipation \citep[e.g][]{2003W,2006F,2007FT,2011NF,2011WL} From the present-day orbits of exoplanets we can potentially distinguish between mechanisms proposed to shape the architectures of planetary systems during the early period of dynamical upheaval. In this spirit, \citet{2011MJa} used the distribution of stellar obliquities to estimate the fraction of hot Jupiters on misaligned orbits and to distinguish between two specific migration mechanisms \citep[see also][]{2009F,2010TC,2010WF}; \citet{2012NF} recently applied a similar technique to estimate the relative contributions of two different mechanisms. However, deducing dynamical histories from the \emph{eccentricity} distribution of exoplanets poses a challenge because most hot Jupiters have already undergone tidal circularization and ``cold'' Jupiters at larger orbital distances may have formed in situ. Furthermore, type-II (gap-opening) migration may either excite or damp a planet's eccentricity through resonance torques \citep{2003G,2004S}. Finally, \citet{2011GR} find evidence that some hot Jupiters may have undergone disk migration either prior to or following scattering. In the latter case, disk migration may have damped their eccentricities. The eccentricity distribution is potentially shaped by a combination of HEM, tidal circularization, and planet-disk interactions. Motivated by the HEM mechanisms proposed by \citet{2003W} and others, S12 proposed an observational test for HEM. As an alternative to modeling the \emph{distribution} of eccentricities, they suggested that we look for the \emph{individual} highly eccentric, long-period progenitors of hot Jupiters caught of the act of tidal circularization. S12 identified HD\,80606\,b as one such progenitor, which was originally discovered by radial velocity (RV) measurements of its host star's reflex motion \citep{2001N} and later found to transit along an orbit that is misaligned with respect to its host star's spin axis \citep{2009M,2009WH}. From statistical arguments S12 predicted that if HEM produces the majority of hot Jupiters, the \kep Mission should detect several ``super-eccentric'' Jupiters with orbital periods less than 93 days and eccentricities in excess of 0.9. A couple of these planets should be proto-hot Jupiters, with post-circularization semimajor axes in the region where all hot Jupiters have circularized (i.e. $P < 5$ days). Several more eccentric planets should have final periods above 5 days, in the region where not all hot Jupiters have circularized; these planets may be ``failed-hot Jupiters'' that will never circularize over their host stars' lifetimes. A failed-hot Jupiter may have either halted at its post-HEM location due to the tidal circularization timescale exceeding the age of the system, or undergone some tidal circularization but subsequently stalled after a perturber in the system raised its periapse. S12's prediction is supported by the existence of super-eccentric eclipsing binaries in the \kep sample, which are also thought to have been created by HEM mechanisms \citep{2012D}. To test the HEM hypothesis we are ``distilling'' eccentric, Jupiter-sized planets from the sample of announced \kep candidates using the publicly released \kep light curves \citep{2011BKB,2012B}. We described the distillation process and our technique for measuring eccentricities from transit light curves based on the ``photoeccentric effect'' in \citet{2012DJ}, hereafter Paper I. In summary, eccentric Jupiters are readily identified by their short ingress/egress/total transit durations \citep{2007B,2008F,2008B,2012P,2012KC}. A Markov-Chain Monte Carlo (MCMC) exploration of the posterior distributions of the transit parameters, together with a loose prior imposed on the stellar density, naturally accounts for the eccentricity-dependent transit probability and marginalizes over the periapse angle, yielding a tight measurement of a large orbital eccentricity (Paper I). Here we present the first eccentric, Jupiter-sized candidate from the \kep sample: \kep Object of Interest (KOI) number 1474.01. We find that this eccentric candidate also has large transit-timing variations (TTVs). In fact, the TTVs are so large that they were likely missed by the automatic TTV-detection algorithms, as they were not listed in a recent cataloging of TTV candidates \citep{2012FR,2012S}. \citet{2011B} recently deduced the presence and planetary nature of the non-transiting Kepler-19c from the TTVs it caused in the transiting planet Kepler-19b, demonstrating the viability of detecting non-transiting planets through TTVs. More recently, \citet{2012N} characterized a Saturn-mass non-transiting planet using this technique. Thus the TTVs of 1474.01 may place constraints on the nature of an additional, unseen companion, thereby elucidating the dynamical history of the system. In \S \ref{sec:lightcurve}, we present the light curve of KOI-1474.01. In \S \ref{sec:star}, we characterize the host-star KOI-1474 using \kep photometry, ground-based spectroscopy, and stellar evolution models. In \S \ref{sec:fpp}, we estimate the candidate's false positive probability (FPP) to be 3.1\%. In \S \ref{sec:ecc}, we measure KOI-1474.01's large eccentricity, investigate its TTVs and the perturbing third body that causes them, and measure the projected alignment of the transiting planet's orbit with the host star's spin axis. In \S \ref{sec:protofail}, we place KOI-1474.01 in the context of known hot Jupiters, proto-hot Jupiters, and failed-hot Jupiters, and explore whether KOI-1474.01 is a failed-hot Jupiter that will retain its current orbit or a proto-hot Jupiter that will eventually circularize at a distance close to the host star. We conclude in \S \ref{sec:discuss} by discussing the implications for planetary system formation models and suggesting directions for future follow up of highly eccentric planets in the \kep sample. \begin{figure}[h] \begin{centering} \includegraphics{fig1.eps} \end{centering} \end{figure} \begin{figure}[t] \begin{centering} \caption{Detrended light curves, color-coded by transit epoch, spaced with arbitrary vertical offsets. The top eight light curves are phased based on a constant, linear ephemeris (Table \ref{tab:planparams}, column 3), revealing the large TTVs. Each light curve is labeled `C' with its best-fit mid-transit time (Table \ref{tab:planparams}, column 3). In the second-from-the-bottom compilation, each light curve is shifted to have an individual best-fitting mid-transit time at t=0. The bottom points are the residuals multiplied by 10. Solid lines: best-fitting eccentric model (Table \ref{tab:planparams}, column 3). \label{fig:lightcurves}} \end{centering} \end{figure} | \label{sec:discuss} We have identified KOI-1474.01 as a highly eccentric, Jupiter-sized planet using a combination of a detailed analysis of the light curve shape and the statistical validation procedure of \citet{2012M}. This makes KOI-1474.01 the second planet or planet candidate with an eccentricity measured solely via the duration aspect of the ``photoeccentric effect,'' joining KOI-686.01 whose eccentricity we measured in Paper I. We measured one component of the angle between the stellar spin axis and the planet's orbit, finding that the degree of misalignment is not currently well-constrained. Based on the variations in KOI-1474.01's transit times, we explored the identity of a perturbing companion; we found the TTVs to be consistent with perturbations from a massive, eccentric outer companion but could not uniquely constrain the perturber's mass, period, eccentricity, and mutual inclination with the currently available data. However, the main reason the perturber's parameters are poorly constrained is that we have only witnessed perturber periapse passage; we are likely to witness another periapse passage over the timespan of the \kep mission, potentially allowing us to distinguish between possible perturbers, including a coplanar giant planet vs. a brown dwarf with a large mutual inclination. Because of the uncertainty in KOI-1474.01's measured orbital eccentricity and possible secular variations in that eccentricity due to the perturbing companion, it is not yet clear whether KOI-1474.01 is a proto-hot Jupiter --- with a periapse close enough to its star that the planet will undergo full tidal circularization over the star's lifetime --- or a failed-hot Jupiter, just outside the reach of fast tidal circularization. However, either way, the planet's discovery adds to the growing evidence that HEM mechanisms play a major role in shaping the architecture of planetary systems. The broad eccentricity distribution of extrasolar planets \citep{2008J}, the sculpting of debris disks by planets on inclined and eccentric orbits \citep[e.g][]{1997ML,1999T,2001A,2006Q,2008L,2009CK,2011D,2012DM}, the population of free-floating planets \citep{2011SK}, and the large mutual inclinations measured in the Upsilon Andromeda system \citep{2010M} all point to a dynamically violent youth for planetary systems. But the strongest evidence for HEM comes from hot Jupiters themselves --- their existence and, in many cases, misaligned or retrograde orbits \citep[e.g][]{2009WJ,2011J,2011T}. As a proto- or failed-hot Jupiter, KOI-1474.01 plays the crucial role of linking hot Jupiters, which are intrinsically rare, to other planetary systems. Even though they make up only a small percentage of the planet population \citep{2010HM,2011HM,2011Y,2012MM,2012W} we focus attention on hot Jupiters because, like meteorites discovered in Antarctica, they are known to come from somewhere else, bringing with them vital information about the past. In contrast, we do not know whether planets at greater orbital distances or of smaller sizes underwent migration, or if they formed in situ \citep[e.g.][]{2009V,2012H}. Moreover, the HEM mechanisms for producing hot Jupiters --- including planet-planet scattering \citep{2011N}, the Kozai mechanism \citep{2003W,2007FT,2011NF}, dynamical relaxation \citep{2008J}, and secular chaos \citep{2011WL} --- make specific predictions for the inclination distributions of hot Jupiters, which can be probed via the Rossiter-McLaughlin effect. The existence of proto- and failed-hot Jupiters will allow us to argue that the mechanisms for producing hot Jupiters are, more generally, the mechanisms that sculpt many types of planetary systems, particularly those with giant planets within ~1 AU. The KOI-1474 system---an inner proto- or failed-hot Jupiter with a massive, long-period companion---may be the prototype of systems of hot Jupiters with distant, massive, outer companions, including as HAT-P-13 \citep{2009BH}, HAT-P-17 (\citealt{2012HB}; a hot Saturn), and Qatar-2 \citep{2012BA}. \citet{2012BA} present a compilation of the eight other hot Jupiters with known outer companions. HD~163607 \citep{2012G} resembles KOI-1474.01 in that it harbors both an eccentric inner planet (e = 0.73, P = 75.29 days) and an outer companion (in this case, a massive outer planet); however, inner planet HD~163607~b is very likely a failed-hot Jupiter, as it has $\pfinal = 24$ days. The expanding baseline for radial-velocity measurements may reveal additional, long-period outer companions of other hot Jupiters, proto-hot Jupiters, and failed-hot Jupiters \citep{2009W}. These additional companions may have been the culprits responsible for the HEM of their inner brethren. Moreover, although \citet{2012SR} examined the transit timing variations of \kep hot Jupiters and found no evidence for \emph{nearby} massive planets, the extended \kep Mission will allow for the detection of distant companions, should they exist, through TTVs. Through radial-velocity follow up with Keck/HIRES we will measure the mass of KOI-1474.01, tighten the measurement of its high eccentricity, place additional constraints on the outer companion, and potentially discover additional bodies in the system. Assuming a Jupiter-like composition to estimate a mass for KOI-1474.01 of $M_p \approx M_{\rm Jup}$, host star KOI-1474 would have an radial velocity semiamplitude of $\sim 70$ m s$^{-1}$, feasible for detection using Keck/HIRES. We will then combine the RV-measured eccentricity with the transit light curves to more tightly constrain the stellar parameters, yielding a better constraint on the planet's line-of-sight spin-orbit angle $|i-i_s|$, which is currently ambiguous due to uncertainty in the stellar radius. It may even be possible to detect the Rossiter-McLaughlin effect, which has a maximum amplitude of $\approx 50$ m/s (\citealt{2010W}, eqn. 40). Although RV measurements of such a faint star ($K_P = 13.005$) pose a challenge, \citet{2012J} have demonstrated the feasibility of following up faint \kep targets with their measurements of KOI-254, a much fainter, redder star ($K_P = 15.979$). KOI-1474.01 contributes to the growing sample of proto- and failed-hot Jupiters. From an estimate of the unbiased number of proto-hot Jupiters, we can determine whether HEM accounts for all the hot Jupiters observed, or whether another mechanism, such as smooth disk migration, must deliver some fraction of hot Jupiters. (See \citealt{2011MJa} for the statistical methodology necessary for such a measurement.) Transiting failed-hot Jupiters orbiting cool stars will be valuable targets for testing the obliquity hypothesis of \citet{2010WF} that hot Jupiters realign cool stars: we would expect failed-hot Jupiters - which have long tidal friction timescales --- to be misaligned around both hot and cool stars. Designed to search for Earth twins in the habitable zones of Sun-like stars, \kep is revealing a wealth of information about the origin of the most unhabitable planets of all: hot Jupiters. \emph{Kepler's} precise photometry, combined with a loose prior on the stellar density, allow us to measure the eccentricities of transiting planets from light curves alone and to search for the highly eccentric proto- and failed-hot Jupiters we would expect from HEM but not from smooth disk migration (S12). If our basic understanding of HEM and tidal circularization is correct, KOI-1474.01 is the first of a collection of highly eccentric planets that will be discovered by \kep. | 12 | 6 | 1206.5579 |
1206 | 1206.2503_arXiv.txt | Supernova theory, numerical and analytic, has made remarkable progress in the past decade. This progress was made possible by more sophisticated simulation tools, especially for neutrino transport, improved microphysics, and deeper insights into the role of hydrodynamic instabilities. Violent, large-scale nonradial mass motions are generic in supernova cores. The neutrino-heating mechanism, aided by nonradial flows, drives explosions, albeit low-energy ones, of ONeMg-core and some Fe-core progenitors. The characteristics of the neutrino emission from new-born neutron stars were revised, new features of the gravitational-wave signals were discovered, our notion of supernova nucleosynthesis was shattered, and our understanding of pulsar kicks and explosion asymmetries was significantly improved. But simulations also suggest that neutrino-powered explosions might not explain the most energetic supernovae and hypernovae, which seem to demand magnetorotational driving. Now that modeling is being advanced from two to three dimensions, more realism, new perspectives, and hopefully answers to long-standing questions are coming into reach. | \label{sec:intro} When, why, and how can the catastrophic infall of the core of a massive star be reversed to trigger the powerful ejection of the stellar mantle and envelope in a supernova (SN) explosion? This fundamental problem of stellar astrophysics has been a matter of intense research since the crucial role of SNe for the synthesis of heavy elements and for the dissemination of the nuclear burning products of stars had been recognized by Burbidge et al.~\cite{Burbidgeetal1957}. The latter authors also noticed that nuclear statistical equilibrium in the hot, dense core of evolved stars (at $T \gtrsim 7\times 10^9$\,K) favors iron dissociation to alpha particles, and they concluded that the huge demand of energy (about 1.7\,MeV per nucleon or $1.7\times 10^{18}\,$erg per gram) must be supplied by gravitational binding energy, leading to a contraction of the stellar core and ultimately to a dynamical implosion on a timescale of less than a second, $t_\mathrm{coll}\sim 0.21/\sqrt{\rho_8}$\,s, when the average density $\rho_8 \equiv \rho/(10^8\mathrm{g/cm^3})$ exceeds unity. This groundbreaking insight is in line with Baade \& Zwicky's earlier idea that SNe could represent the transition of ordinary stars to neutron stars (NSs)~\cite{BaadeZwicky1934}. Already in 1960 Hoyle \& Fowler~\cite{HoyleFowler1960} proposed the two basic scenarios of stellar death: thermonuclear runaway at degenerate conditions (which, as we know now, drives the destruction of white dwarf stars in Type~Ia SNe) and the implosion of stellar cores (associated with what is called core-collapse supernovae (CCSNe) of Types II, Ib/c, and hypernovae\footnote{Observationally, SNe~II exhibit strong H-Balmer lines in their early spectra, whereas SNe~I show no H-lines. In SNe~Ia there are Si-lines, in SNe~Ib no Si- but He-lines, and in SNe~Ic none of these, indicating explosions of stars that had lost their hydrogen envelope or both the outer hydrogen and helium shells before collapse. More sub-classes have been introduced, some of them motivated only by recent discoveries: SNe~II-P and II-L are discriminated by a plateau phase or linear decay of their lightcurves after the peak, IIb events have only thin H-shells left, and spectra of IIa and IIn cases possess signatures of a dense circumstellar medium.}). They hypothesized (following~\cite{Burbidgeetal1957}) that the gravitational compression of the core raises the temperature such that thermonuclear fuel could be ignited to release the energy for triggering the ejection of the outer parts of the star. They also mentioned simulations by Colgate \& Johnson~\cite{ColgateJohnson1960,ColgateGrasbergerWhite1961}, in which the ``bounce'' of a forming NS launched a spherical shock wave that reversed the infall of the overlying stellar shells to make them gravitationally unbound. Colgate \& White~\cite{ColgateWhite1966} realized that gravitational binding energy of order $E_\mathrm{b}\sim GM_\mathrm{ns}^2/R_\mathrm{ns} > 10^{53}$\,erg, which is released when the core of a star collapses to a NS, is converted to neutrino emission and provides a huge energy reservoir for powering the SN blast wave. They argued correctly that in stellar layers pulled inward at supersonic speed along with the imploding core, thermonuclear combustion is unable to initiate an outward acceleration. Instead they proposed that a fraction of the intense neutrino flux may get absorbed in the mantle of the star to cause the explosion. More than four decades of theoretical and numerical modeling work, spearheaded by early pioneers of the field like Dave Arnett, Jim Wilson, Hans Bethe, Gerry Brown, Steve Bruenn, Wolfgang Hillebrandt, Jim Lattimer, and David Schramm, have helped to sharpen our picture of the diverse physical ingredients and processes that play a role in the core of dying stars, among them magnetohydrodynamic (MHD) effects, fluid instabilities and turbulent flows, the finite-temperature equation of state (EoS) of NS matter, neutrino transport and neutrino-matter interactions, and general relativistic gravity. While the bounce-shock mechanism is not supported by any modern simulation with state-of-the-art treatment of the physics, the ``delayed neutrino-heating mechanism'' as discussed by Bethe and Wilson~\cite{BetheWilson1985} and aided by violent, nonradial mass motions in the collapsing stellar core~\cite{Herantetal1994,Burrowsetal1995,JankaMueller1995,JankaMueller1996}, has advanced to the widely favored scenario for powering the majority of SNe. The momentum behind the quest for solving the puzzle of the SN mechanism originates from important questions at the interface of astrophysics and nuclear, particle, and gravitational physics, for example: \begin{itemize} \item What is the link between the properties of SNe and their progenitor stars? \item Which stars collapse to black holes (BHs) instead of NSs, which fraction of stellar collapses do not yield explosions? \item What are the birth properties of the compact remnants, i.e.\ their masses, spins, magnetic fields, and recoil velocities? \item How can the high velocities of young pulsars be explained? Is any exotic physics necessary? \item What characteristics does the neutrino burst from a SN have and what does it tell us about neutrino properties and the extreme conditions in the newly formed NS? \item What is the gravitational-wave signature of a stellar collapse event and which information can we extract about the dynamical processes in the SN core? \item What is the nucleosynthetic role of massive star explosions in the chemogalactic history? \item Are SNe the long-sought sources of r-process elements, in particular also of the lanthanides, the third abundance peak, and actinides? \item What is the population-integrated energetic footprint left by SN explosions in the dynamical evolution of galaxies? \end{itemize} In the following sections we will review the known types of stellar collapse events (Sect.~\ref{sec:progenitors}), the ingredients and current status of numerical modeling (Sect.~\ref{sec:numericalmodeling}), the mechanisms by which massive stars might explode (Sect.~\ref{sec:mechanisms}), and the signatures of the explosion mechanism that might serve for observational diagnostics (Sects.~\ref{sec:coresignals} and \ref{sec:remnants}). We will provide an update of recent developments as follow-up and supplement of previous reports that have approached the topic from different perspectives~\cite{Fryer2004,Mezzacappa2005,WoosleyJanka2005,WoosleyBloom2006,Kotakeetal2006,Jankaetal2007,Ott2009,Thielemannetal2011}. \begin{figure}% \centerline{\psfig{figure=JankaFig1.eps,width=12truecm}} \caption{{\small Stellar death regions with schematic stellar evolution tracks in the plane of central density ($\rho_\mathrm{c}$) and central temperature ($T_\mathrm{c}$). Colored death regions are labeled by the instability process causing the collapse of the stellar core, and the blue tracks are labeled by the corresponding rough birth-mass range of objects reaching the different stages of central burning (indicated by red dashed lines). Yellow diagonal lines mark the beginning of degeneracy (short-dashed) and strong degeneracy (long-dashed) of the electron plasma. Note that realistic stellar tracks exhibit wiggles and loops when the ignition of the next burning stage is reached and the stellar core adjusts to the new energy source (see Ref.~\cite{Wheeler1990}.)}} \label{jankafig1} \end{figure} \begin{figure}% \centerline{\psfig{figure=JankaFig2.eps,width=10truecm}} \caption{{\small Core-density profiles of different SN progenitors at the onset of gravitational collapse. The black line corresponds to the ONeMg core of an 8.8\,$M_\odot$ star~\cite{Nomoto1984}, the other three are SN progenitors with iron cores: an 8.1\,$M_\odot$ ultra metal-poor ($10^{-4}$ solar metallicity) star (A.~Heger, private communication) and 11.2\,$M_\odot$~\cite{Woosleyetal2002} and 15\,$M_\odot$~\cite{WoosleyWeaver1995} solar-metallicity stars. The steps and kinks in the curves correspond to composition-shell interfaces (Fe/Si and O/C for the 11.2 and 15\,$M_\odot$ models and inner and outer boundaries of a C-O-Ne-layer for the 8.1\,$M_\odot$ case).}} \label{jankafig2} \end{figure} | Supernova theory has made remarkable progress over the past decade, promoted by common interests of the astro-, particle (neutrino), nuclear, and gravitational physics communities and by an increasing number of active (young) researchers in the field. A deeper understanding of the physical mechanisms that initiate and fuel SN and HN explosions of massive stars is of crucial importance not only for establishing the progenitor-remnant connection but also for predicting the properties of stellar explosions, their nucleosynthetic output, and the characteristics of their gravitational-wave and neutrino signals. The most sophisticated present simulations demonstate that neutrino-energy deposition can power ECSNe (even in spherical models) of $\sim$9\,$M_\odot$ stars with ONeMg-cores near the lower mass limit for SN progenitors (Fig.~\ref{jankafig5}). Overall, the features of such explosions, e.g., low energy and little nickel production, seem to be compatible with observational candidates like the Crab SN and some faint transients. Multi-dimensional simulations suggest these explosions to be potential sources of light r-process nuclei up to silver and palladium (Sect.~\ref{sec:heavyelements}). Several groups have also reported successful neutrino-driven explosions (with multi-group neutrino transport) for Fe-core progenitors above 10\,$M_\odot$ (Sect.~\ref{sec:hydromodels}; Figs.~\ref{jankafig4}, \ref{jankafig5}). Ultimate confirmation of the viability of this mechanism for a wider range of progenitor masses therefore seems to be in reach. The onset of the explosion can be understood as a global runaway instability of the accretion layer, whose initiation depends on the power of neutrino-energy deposition. While the exact mode of the runaway is still a matter of exploration and debate (e.g., low-multipole SASI or higher-multipole convective, oscillatory or nonoscillatory?), its threshold in terms of the driving neutrino luminosity is lowered by nonradial fluid motions in the neutrino-heating layer. Such flows play a supportive role because they stretch the residence time of matter in the gain region and thus decrease the heating timescale and increase the efficiency of neutrino-energy deposition, leading to successful explosions even when sophisticated spherical models fail (Sect.~\ref{sec:nonsphericaleffects}). The efficiency of neutrino-energy transfer, the growth conditions and growth rates of different hydrodynamic instabilities, and the critical luminosity threshold for an explosion may not only depend on the dimension and thus will ultimately require simulations in 3D, but have been shown to depend also on putative ``details'' of the physics ingredients like approximations for the energy and velocity dependence of the neutrino transport, the neutrino-interaction rates, general relativity, and the contraction of the nascent NS in response to the nuclear EoS (Sect.~\ref{sec:numericalmodeling}; Fig.~\ref{jankafig4}). Moreover, the outcome of the complex neutrino-hydrodynamical simulations can be sensitive to the numerical resolution, which naturally is subject to limitations in full-scale, multi-dimensional SN-core models. While detailed modeling of the processes in collapsing stars now pushes forward from the second to the third dimension, facing considerable computational challenges and demands mainly for the neutrino transport, a growing host of studies begins to explore the observational consequences of neutrino-driven explosions. In view of existing and upcoming big detection facilities, in particular neutrino and gravitational-wave signals (Figs.~\ref{jankafig5}, \ref{jankafig7}, \ref{jankafig8}) are of relevance for SN-core diagnostics targeting a future Galactic SN. The former even have the potential to yield valuable information on particle properties of the neutrinos provided the characteristics of the SN emission are sufficiently well understood (e.g., \cite{Kachelriessetal2005,Chakrabortyetal2011C,Ellisetal2011,Ellisetal2012}). Sophisticated neutrino transport and interaction treatments have revealed interesting signal features like an amazing robustness of the neutronization $\nu_e$ burst~\cite{Kachelriessetal2005}, characteristic differences of the rise time of the $\bar\nu_e$ and $\nu_x$ emission after bounce~\cite{Chakrabortyetal2011C}, luminosity variations associated with nonsteady flows in the accretion layer~\cite{MarekJankaMueller2009,Brandtetal2011}, and a close similarity of the luminosities and spectra of neutrinos and antineutrinos of all flavors during the PNS cooling phase (Sect.~\ref{sec:neutrinos}, Fig.~\ref{jankafig5}; \cite{Huedepohletal2010,Fischeretal2011C}) with important consequences for SN nucleosynthesis (Sect.~\ref{sec:heavyelements}). While a Galactic SN in the near future is a realistic possibility, it will be a unique event and might not provide evidence of wider validity. Photometric and spectroscopic diagnostics of extragalactic SNe and of gaseous, young SN remnants, which reveal information on explosion energies, $^{56}$Ni production, ejecta masses, asymmetries, and composition, as well as progenitor constraints (cf.\ Fig.~\ref{jankafig3}) are therefore extremely valuable, and more is desirable. First-principle explosion models begin to become mature enough to be linked to such observations, a possibility that defines a fruitful territory for future reseach. Neutrino-driven explosion models also begin to allow for predictions of compact remnant (NSs and BHs) masses, kicks, and spins. Nonradial hydrodynamic instabilities in the collapsing stellar core, which can grow from small, random initial perturbations before neutrino heating revives the stalled shock, lead to low-multipole asymmetries that trigger anisotropic and inhomogeneous expulsion of matter. Hydrodynamic instabilities in the SN core therefore do not only yield a natural explanation of the origin of pulsar kicks up to more than 1000\,km/s (Sect.~\ref{sec:pulsarkicks}, Fig.~\ref{jankafig9}); they also seed large-scale mixing processes in the exploding star, accounting for the penetration of high-velocity clumps of inner-core material into the hydrogen and helium ejecta of well observed SN explosions (Sect.~\ref{sec:snasymmetries}, Fig.~\ref{jankafig10}). First results of a systematic exploration of the progenitor-supernova connection based on the neutrino-heating mechanism show strong sensitivity of the explosion properties on the stellar structure and, for the employed set of stellar models~\cite{Woosleyetal2002}, large variations even within narrow progenitor-mass intervals (Sect.~\ref{sec:NSBH}; Fig.~\ref{jankafig11}). The explosion models can reproduce fundamental properties of the empirical remnant-mass distribution but reveal that neutrino-driven explosions are unlikely to explain SN energies above $\sim 2\times 10^{51}$\,erg and nickel masses significantly higher than 0.1\,$M_\odot$. This underlines the need for an alternative engine that powers stellar blast waves with energies from several $10^{51}$\,erg up to more than $10^{52}$\,erg. Such hyperenergetic events, which typically also exhibit unusually large nickel ejection (Fig.~\ref{jankafig3}) and deformation, are most probably energized by magnetorotational effects. Many questions remain to be answered in this context and require more observations and theoretical work. What discriminates progenitors of ``normal'' SNe from those of HNe? Is rapid rotation of the progenitors the crucial parameter? Is it connected to binary evolution? Is there a continuous spectrum of stellar explosions connecting the SN and HN regimes? Is a mixed mechanism, neutrino-heating in combination with magnetorotational energy transfer, at work in such events? On the theory side the mission of clarifying the SN engines is severely handicapped by the unavailability of multi-dimensional stellar evolution models with the quality to reduce the major uncertainties of the stellar structure, rotation, and magnetic fields at the onset of core collapse. It is clear that reliable theoretical predictions of the progenitor-remnant connection and of explosion properties ---energies, nucleosynthetic yields, asymmetries, remnant masses, and neutrino and GW signals--- heavily depend on a firm knowledge of the stellar conditions at the time the gravitational instability is reached. | 12 | 6 | 1206.2503 |
1206 | 1206.2029_arXiv.txt | We present $R$-Band light curves of Type~II supernovae (SNe) from the Caltech Core Collapse Project (CCCP). With the exception of interacting (Type~IIn) SNe and rare events with long rise times, we find that most light curve shapes belong to one of three distinct classes: plateau, slowly declining and rapidly declining events. The last class is composed solely of Type~IIb SNe which present similar light curve shapes to those of SNe Ib, suggesting, perhaps, similar progenitor channels. We do not find any intermediate light curves, implying that these subclasses are unlikely to reflect variance of continuous parameters, but rather might result from physically distinct progenitor systems, strengthening the suggestion of a binary origin for at least some stripped SNe. We find a large plateau luminosity range for SNe~IIP, while the plateau lengths seem rather uniform at approximately $100$ days. As analysis of additional CCCP data goes on and larger samples are collected, demographic studies of core collapse SNe will likely continue to provide new constraints on progenitor scenarios. | Type~II supernovae (SNe) are widely recognized as the end stages of massive H-rich stars and represent the bulk of observed core collapse SNe (see Filippenko 1997 for a review of SN classifications). Several sub-types of Type~II SNe have been observed. Those showing a plateau in their light curve are known as Type~IIP events, while those showing a linear decline from peak magnitude are classified as IIL. A third class of events, Type~IIb, characterized by its spectral rather than its photometric properties, develops prominent He features at late times. Finally, Type~IIn SNe display narrow lines in their spectra, indicative of interaction between the SN ejecta and a dense circum stellar medium. Red supergiants (RSGs) have been directly identified as the progenitors of Type~IIP SNe for SN2003gd (Van Dyk et al. 2003; Smartt et al. 2004), SN2004A (Hendry et al. 2006), SN2005cs (Maund et al. 2005; Li et al. 2006), SN2008bk (Mattila et al. 2008; Van Dyk et al. 2012) and SN2009md (Fraser et al. 2011); see Smartt 2009 for a review. Such stars have thick hydrogen envelopes that are ionized by the explosion shock wave. As the shocked envelope expands and cools, it recombines, releasing radiation at a roughly constant rate, thus producing a plateau in the light curve (e.g. Popov 1993; Kasen \& Woosley 2009). It follows that SNe IIL might be the explosions of stars with less massive H envelopes that can not support a plateau in their light curve. SN IIb progenitors, then, would contain an even smaller H mass. However, if SNe IIP-IIL-IIb progenitors represent merely a sequence of decreasing H-envelope mass, one would expect the properties of these SNe to behave as a continuum. Specifically, a gradual transition in light curve shape should be observed when examining a homogeneous sample of events. The Caltech Core Collapse Project (CCCP; Gal-yam et al. 2007) is a large observational survey which made use of the robotic 60-inch (P60; Cenko et al. 2006) and Hale 200-inch telescopes at Palomar Observatory to obtain optical $BVRI$ photometry and spectroscopy of 48 nearby core collapse SNe. By providing a fair sample of core collapse events with well defined selection criteria and uniform, high quality optical observations, CCCP allows to study core collapse SNe as a population rather than as individual events. Light curves of Type~Ib/c SNe from CCCP have been presented and analyzed by Drout et al. (2011). Type~IIn CCCP events are treated by Kiewe et al. (2012). Here we present photometry of 21 non-interacting Type~II SNe with well observed light curves collected through CCCP. We present $R$-Band data for most of the events to simplify the comparison of their light curve shapes. A more detailed multi-color analysis will be presented in a forthcoming paper. | We plot the $R$-Band light curves of $15$ Type~II events normalized to peak magnitude in Figure \ref{alllcs} (top panel). Rather than forming a continuum, we find that the light curves group into three distinct sub-classes: plateau, slowly declining (1-2 Mag/100 days) and initially rapidly declining (5-6 Mag/100 days) events (see also Table \ref{subtypes}). We note that the three rapidly declining events are all Type~IIb and that they display similar light curve shapes to those of Type~Ib/c SNe (Drout et al. 2011). We perform a Kolmogorov-Smirnov test and find that the probability that the measured $\Delta\textrm{M}15_{R}$ values for the rapid and slow decline groups are drawn from a single underlying distribution is $2\%$. Three events (SN2004ek, SN2005ci and SN2005dp; Figure \ref{pecs}) display prolonged rising periods in their light curves. They do not show signs of interaction in their spectra and may be explosions of compact blue supergiant progenitors (Kleiser et al. 2011; Pastorello et al. 2012), as demonstrated directly in the case of SN 1987A (see Arnett et al. 1989 for a review). Finally, one event (SN2004em; Figure \ref{pecs}) displays a very peculiar photometric behavior. For the first few weeks it is similar to a Type~IIP SN, while around day $25$ it suddenly changes behavior to resemble a SN 1987A-like event. The full photometric dataset is available online through WISeREP\footnote{http://www.weizmann.ac.il/astrophysics/wiserep} (Yaron \& Gal-Yam 2012). \begin{figure*} \includegraphics[width=18cm]{Fig1a.eps} \includegraphics[width=18cm]{Fig1b.eps} \caption{Top Panel: $R$-band light curves of 15 Type~II SNe from CCCP (excluding the four events presented in Figure \ref{pecs}), normalized in peak magnitude (SN2004fx data taken from Hamuy et al. 2006; SN2005ay data taken from Gal-Yam et al. 2008b; SN2005cs data taken from Pastorello et al. 2009). A clear subdivision into three distinct subtypes is apparent: plateau, slowly declining and rapidly declining SNe (the latter consisting only of SNe IIb). Reference SNe are shown for comparison (SN1999em from Leonard et al. 2002; SN2009kr from Fraser et al. 2010, found to be a member of the IIL subclass as claimed by Elias-Rosa et al. 2010; SN1993J from Richmond et al. 1994; SN2011dh from Arcavi et al. 2011). We also overplot the SN Ib/c template derived by Drout et al. (2011). The data have been interpolated with spline fits (except for SN2005by, where a polynomial fit provided a better trace to the data). The shaded regions denote the average light curve $\pm2\sigma$ of each subclass. The maximal 7-day uncertainty in determining the explosion times is illustrated by the interval in the top right corner. Bottom Panel: $R$-band light curves of 9 Type~IIP SNe from CCCP with respect to their estimated explosion time (except for SN2005au and SN2005bw, marked by dashed lines, for which the explosion date is not known to good accuracy). The light curve of SN2004fx is taken from Hamuy et al. (2006), that of SN2005ay from Gal-Yam et al. (2008b) and that of SN2005cs from Pastorello et al. (2009). SN1999em (Leonard et al. 2002) is shown for comparison. A spread in plateau luminosities is apparent while plateau lengths seem to converge around 100 days. Spline fits were applied to the data.} \label{alllcs} \end{figure*} \begin{figure*} \includegraphics[width=18cm]{Fig2.eps} \caption{$R$-Band photometric data of the CCCP events included in Figure \ref{alllcs} together with the spline fits shown in that figure. The x-axis is MJD-53000 in days, and the y-axis is apparent magnitude. The light curve of SN2004fx is taken from Hamuy et al. (2006), that of SN2005ay from Gal-Yam et al. (2008b) and that of SN2005cs from Pastorello et al. (2009).} \label{spline} \end{figure*} \begin{deluxetable*}{llllllll} \tablecolumns{8} \tablewidth{0pt} \tablecaption{Subdivision of the CCCP Type~II SN light curves analyzed. For the declining SNe, the peak magnitudes and the $\Delta\textrm{M}15_{R}$ parameter (denoting the magnitude drop at 15 days after peak in the $R$-band) are derived from the smooth fits to the light curves (limits are stated when the rise to peak is not detected in the data). For the plateau events, the average luminosity during the first 50 days is taken as the plateau magnitude (one $\sigma$ shown in parentheses). The explosion dates assumed in Fig. \ref{alllcs} are noted (explosion date window widths shown in parentheses).} \tablehead{ \colhead{SN} & \colhead{Type} & \colhead{Peak/Plateau Mag} & \colhead{$\Delta\textrm{M}15_{R}$} & \colhead{DM} & \colhead{DM Source} & \colhead{Explosion MJD} & \colhead{Reference} \\ } \startdata SN2004du & Plateau & $-17.464$ ($0.055$) & & $34.26$ & NED & $53229$ ($2$) & IAUC 8387 \\ SN2004er & Plateau & $-16.669$ ($0.077$) & & $33.39$ & NED & $53273$ ($2$) & IAUC 8412 \\ SN2004et & Plateau & $-17.482$ ($0.088$) & & $28.80$ & NED & $53268$ ($4$) & IAUC 8413 \\ SN2004fx\footnote{photometry from Hamuy et al. 2006} & Plateau & $-16.087$ ($0.038$) & & $32.60$ & NED & $53300$ ($3$) & IAUC 8431 \\ SN2005aa & Plateau & $-15.824$ ($0.134$) & & $34.83$ & Redshift & $53402$ ($5$) & IAUC 8476 \\ SN2005au & Plateau & $-17.069$ ($0.141$) & & $34.07$ & NED & & IAUC 8496 \\ SN2005ay\footnote{photometry from Gal-Yam et al. 2008b} & Plateau & $-16.447$ ($0.119$) & & $31.21$ & NED & $53452$ ($4$) & IAUC 8500/2 \\ SN2005bw & Plateau & $-16.945$ ($0.150$) & & $35.15$ & Redshift & & CBET 147 \\ SN2005cs\footnote{photometry from Pastorello et al. 2009} & Plateau & $-15.040$ ($0.079$) & & $29.62$ & V12 & $53548$ ($2$) & IAUC 8553\\ \hline SN2005Z & Slow Decline & $<-17.5$ & $0.135$ & $34.61$ & Redshift & $53391$ ($5$) & IAUC 8476 \\ SN2005ab & Slow Decline & $<-15.2$ & $0.405$ & $34.14$ & Redshift & $53406$ ($7$) & IAUC 8478 \\ SN2005an & Slow Decline & $<-17.0$ & $0.170$ & $33.39$ & Redshift & $53430$ ($7$) & CBET 113 \\ SN2005ba & Slow Decline & $-17.4$ & $0.202$ & $35.60$ & Redshift & $53450$ ($7$) & IAUC 8503 \\ Quest SN1 & Slow Decline & $-16.5$ & $0.181$ & $33.63$ & Redshift & $53256$ ($7$) & \\ \hline SN2004ex & Rapid Decline & $-17.2$ & $0.751$ & $34.41$ & Redshift & $53289$ ($1$) & IAUC 8418 \\ SN2005bp & Rapid Decline & $<-16.8$ & $0.926$ & $35.41$ & Redshift & $53477$ ($7$) & IAUC 8515 \\ SN2005by & Rapid Decline & $-17.5$ & $0.750$ & $35.39$ & Redshift & $53482$ ($6$) & IAUC 8523 \\ \hline SN2004ek & Prolonged Rise & $-18.2$ & & $34.38$ & NED & & IAUC 8405 \\ SN2005ci & Prolonged Rise & $-16.1$ & & $32.80$ & NED & & IAUC 8541 \\ SN2005dp & Prolonged Rise & $-17.6$ & & $32.57$ & NED & & IAUC 8591\\ SN2004em & Peculiar & $-17.9$ & & $34.05$ & Redshift & & IAUC 8406 \enddata \label{subtypes} \end{deluxetable*} \subsection{Declining SNe} Aside from establishing a different rate of decline for SNe IIb compared to SNe IIL, Figure \ref{alllcs} (top panel) suggests that the IIP, IIL and IIb subtypes do not span a continuum of physical parameters, such as H envelope mass. Rather, additional factors should be considered. Specifically, Type~IIb events might arise from binary systems (as suggested also by recent progenitor studies for SN1993J, Maund et al. 2004; SN2008ax, Crockett et al. 2008; SN2011dh, Arcavi et al. 2011, Van-Dyk et al. 2011). The similarity of the Type~IIb light curves to those of Type~Ib events (also seen in the Drout et al. 2011 data), in addition to the known spectral similarities at late times and the similar peak radio luminosities (Chevalier \& Soderberg 2010), suggests that these two types of events might come from similar progenitor systems. \subsection{Plateau SNe} The $R$-Band light curves of the Type~IIP SNe, on an absolute magnitude scale, can be seen in Figure \ref{alllcs} (bottom panel). We find a wide range of plateau luminosities, but do not have enough statistics to test whether they form a continuous distribution or if there are two distinct underlying types (bright and faint), as previously suggested (Pastorello et al. 2004). The plateau lengths, however, seem rather uniform at $\sim100$ days (with the sole exception of SN2004fx, displaying a shorter plateau)\footnote{Note that SN2005au and SN2005bw are plotted only to show their plateau luminosity, their plateau lengths are unknown due the lack of sufficient constraints on their explosion time}. This is consistent with the plateau length scale given by Popov (1993), which assumes a constant opacity: $$t_{p}=99\frac{\kappa_{0.34}^{1/6}M_{10}^{1/2}R_{0,500}^{1/6}}{E_{51}^{1/6}T_{\textrm{ion},5054}^{2/3}}\textrm{days}$$ (where the radius and mass are expressed in solar units). However, according to this scaling, one would expect to see also longer plateaus (up to $\sim130$ days for $17M_\odot$ progenitors). Kasen \& Woosley (2009) show that heating from radioactive decay of $^{56}\textrm{Ni}$ should further extend the duration of the plateau ($0.1M_\odot$ of $^{56}\textrm{Ni}$ extends the plateau by $\sim24\%$ in their models, without greatly affecting the plateau luminosity). We do not find any of these long duration plateaus in our sample, though a few such events have been observed (e.g. Hamuy 2002). The scarcity of observed short plateaus and the possibly related sharp distinction between IIP and IIL light curves is evident. Such a pronounced absence of intermediate events might suggest that Type~IIL SNe are powered by a different mechanism than that associated with SNe~IIP (e.g., magnetars; Kasen \& Bildsten 2010). \begin{figure*} \includegraphics[width=18cm]{Fig3.eps} \caption{$BVRI$ light curves of the CCCP events not included in Figure \ref{alllcs}. Three events (SN2004ek, SN2005ci, SN2005dp) show long rise times (possibly associated with blue supergiant explosions; Kleiser et al. 2011; Pastorello et al. 2012), while one peculiar event (SN2004em) changes behavior from flat to rising around three weeks after explosion.} \label{pecs} \end{figure*} | 12 | 6 | 1206.2029 |
1206 | 1206.6116_arXiv.txt | We analyze the spectral energy distributions (SEDs) of Lyman break galaxies (LBGs) at \uvdrops\ selected using the \emph{Hubble Space Telescope} (\emph{HST}) Wide Field Camera 3 (WFC3) UVIS channel filters. These \emph{HST}/WFC3 observations cover about 50~\sqmin\ in the GOODS-South field as a part of the WFC3 Early Release Science program. These LBGs at \uvdrops\ are selected using dropout selection criteria similar to high redshift LBGs. The deep multi-band photometry in this field is used to identify best-fit SED models, from which we infer the following results: (1) the photometric redshift estimate of these dropout selected LBGs is accurate to within few percent; (2) the UV spectral slope $\beta$ is redder than at high redshift ($z>3$), where LBGs are less dusty; (3) on average, LBGs at \uvdrops\ are massive, dustier and more highly star-forming, compared to LBGs at higher redshifts with similar luminosities ($0.1L^*\lesssim L \lesssim 2.5L^*$), though their median values are similar within 1$\sigma$ uncertainties. This could imply that identical dropout selection technique, at all redshifts, find physically similar galaxies; and (4) the stellar masses of these LBGs are directly proportional to their UV luminosities with a logarithmic slope of $\sim$0.46, and star-formation rates are proportional to their stellar masses with a logarithmic slope of $\sim$0.90. These relations hold true --- within luminosities probed in this study --- for LBGs from $z\simeq 1.5$ to $5$. The star-forming galaxies selected using other color-based techniques show similar correlations at $z\simeq 2$, but to avoid any selection biases, and for direct comparison with LBGs at $z>3$, a true Lyman break selection at $z\simeq 2$ is essential. The future \emph{HST} UV surveys, both wider and deeper, covering a large luminosity range are important to better understand LBG properties, and their evolution. | The high redshift frontier has moved to $z>7$ as a result of the high resolution near-infrared (NIR) images from the \emph{Hubble Space Telescope} (\emph{HST}) Wide Field Camera 3 (WFC3), and the Lyman break `dropout' technique. The Lyman break technique was first applied to select Lyman break galaxies (LBGs) at $z\simeq 3$ \citep{guha90,stei96,stei99}, and since then it has been extensively used to select and study LBG candidates at redshifts $z\simeq 3$--$8$ \citep[e.g.,][]{bouw07,hath08b,redd09,fink10,yan10}. This dropout technique has generated large samples of faint star-forming galaxy candidates at $z\simeq 3$--$8$. However, at highest redshifts ($z>3$), it is very difficult to understand the details of their stellar populations using current space and ground-based telescopes. Their faint magnitudes make it extremely difficult to do spectroscopic studies, and limited high resolution rest-frame optical photometry make it challenging to investigate their spectral energy distributions (SEDs). These limitations make it imperative to identify and study LBGs at lower redshifts ($z\lesssim 3$). The primary reason for the lack of dropout selected LBGs at \uvdrops\ is that we need highly sensitive space-based cameras to observe the mid- to near-ultraviolet (UV) wavelengths required to identify Lyman break at \uvdrops. The peak epoch of global star-formation rate at \uvdrops\ is now accessible using the dropout technique with the WFC3 UVIS channel. \citet[][hereafter H10]{hath10} and \citet{oesc10} have used the \emph{HST} WFC3 with its superior sensitivity to photometrically identify lower redshift (\uvdrops) LBGs. Understanding the LBGs at $z\lesssim 3$ is vital for two main reasons. First, we need to study the star-formation properties of these LBGs, because they are at redshifts corresponding to the peak epoch of the global star-formation rate \citep[e.g.,][]{redd08,redd09,ly09,bouw10}. Second, they are likely lower redshift counterparts of the high redshift LBGs --- because of their identical dropout selection and similar physical properties --- whose understanding will help shed light on the process of reionization in the early universe \citep[e.g.,][]{labb10,star10}. There are primarily three techniques to select star-forming galaxies at $z\simeq 2$: (1) \emph{sBzK} \citep[using the $B$, $z$, $K$ bands,][]{dadd04,dadd07}, (2) BX/BM \citep[using the $U$, $G$, $R$ bands,][]{stei04, adel04}, and (3) LBG \citep[using the bands which bracket the redshifted Lyman limit, H10;][]{oesc10}. All these approaches select star-forming galaxies, and yield insight into the star-forming properties of these galaxies, but they have differing selection biases, and so these samples don't completely overlap \citep[see][for details]{ly11,habe12}. Therefore, it is essential to apply identical selection criteria at all redshifts to properly compare galaxy samples and accurately trace their evolution. The LBG selection is widely used to select high redshift ($z>3$) galaxies, and to do equal comparison with these galaxies, here we investigate physical properties of LBGs at $z\lesssim 3$. H10 used UV observations of the WFC3 Science Oversight Committee Early Release Science extragalactic program (PID: 11359, PI: O'Connell; hereafter ``ERS''), which covers approximately 50~\sqmin\ in the northern-most part of the Great Observatories Origins Deep Survey \citep[GOODS;][]{giav04} South field, to identify LBGs at \uvdrops. The high sensitivity of the WFC3 UVIS channel data \citep{wind11}, along with existing deep optical data obtained with the Advanced Camera for Surveys (ACS) as part of the GOODS program are ideal to apply dropout technique in observed UV filters to select LBG candidates at \uvdrops. In this paper, we use this H10 sample of LBGs to investigate their physical properties by fitting stellar synthesis models to their observed SEDs. This paper is organized as follows: In \secref{data}, we summarize the WFC3 ERS observations, and discuss our LBG sample at \uvdrops\ as well as the comparison sample of LBGs at \bvdrops. In \secref{seds}, we fit observed SEDs of LBGs at \uvdrops\ and \bvdrops\ to stellar population synthesis models, and discuss the best-fit parameters (redshift, UV spectral slope, stellar mass, stellar age, and star-formation rates) obtained from these SED fits. In \secref{results}, we discuss correlations between best-fit physical parameters and their implications on our understanding of LBGs. In \secref{conclusion}, we conclude with a summary of our results. In the remaining sections of this paper we refer to the \emph{HST}/WFC3 F225W, F275W, F336W, F098M, F125W, F160W, filters as \wfcfuv, \wfcnuv, \wfcuv, \wfcy, \wfcj, \wfch, to the \emph{HST}/ACS F435W, F606W, F775W, F850LP filters as \acsb, \acsv, \acsi, \acsz, and to the \emph{Spitzer}/IRAC 3.6~$\mu$m, 4.5~$\mu$m filters as [3.6], [4.5], respectively, for convenience. We assume a \emph{Wilkinson Microwave Anisotropy Probe} (WMAP) cosmology with $\Omega_m$=0.274, $\Omega_{\Lambda}$=0.726 and \Ho=70.5~km s$^{-1}$ Mpc$^{-1}$, in accord with the 5 year WMAP estimates of \citet{koma09}. This corresponds to a look-back time of 10.4~Gyr at $z\simeq 2$. Magnitudes are given in the AB$_{\nu}$ system \citep{oke83}. | \label{results} \subsection{Stellar Mass vs UV Luminosity Relation}\label{mass_lum} The rest-frame UV light traces recent or instantaneous SFR, while rest-frame optical and NIR data help us to estimate stellar masses of galaxies. If the galaxy stellar mass and UV luminosity are related then we can directly use rest-frame UV light to estimate stellar mass without needing rest-frame optical/NIR data. \figref{fig:mass_abs} shows stellar mass of LBGs at $z\simeq 1.5$--$5$ as a function of their UV absolute magnitude. These quantities are based on best-fit SEDs, and their typical uncertainties are shown in the lower-left corner. The dotted lines are best-fit line obtained by keeping the logarithmic slope fixed at 0.46, which was estimated by \citet{sawi12} for star-forming galaxies at $z\simeq 2$. The dot-dash lines show the scatter from the best-fit line, which is $\sim$0.3 dex for LBGs at \uvdrops\ and about 0.2 dex for LBGs at \bvdrops. We also tested the validity of this relation by fitting the slope of the line rather than fixing it. We find that the fitted slope is in the range of 0.42$\pm$0.06 for our LBG samples, which is consistent with 0.46 within the estimated 1$\sigma$ scatter in this relation. Therefore, we find that a proportionality relation between these two parameters with a logarithmic slope of 0.46 provides a good fit to the data. The stellar masses of the brighter LBGs --- with UV luminosities near the $L_{\rm uv}^*$ value of LBGs at $z\simeq 3$ from \citet{stei99} --- are about a factor of 2 lower than 10$^{10}$~M$_{\odot}$ estimated by \citet{papo01}. This discrepancy, though within our estimated uncertainties, could be due to the fact that we include emission lines in our SED fitting which could affect stellar masses by as much as a factor of $\sim$2. The stellar mass--UV luminosity relation is fairly tight with a small scatter ($\lesssim\,$0.3 dex), which is consistent with other studies at similar redshifts \citep[e.g.,][]{papo01}, and it points to a nearly constant mass-to-light ratio (log(M/L)$\,\simeq\,$--0.5) for LBGs between $z\simeq 1.5$ and $5$. The tightness/lower scatter of the stellar mass--UV luminosity relation in \figref{fig:mass_abs} could be --- in part --- due to the fact that both these quantities are output parameters from the best-fit SEDs, and therefore, it is possible that these parameters are not totally independent. We have addressed this issue and discussed its implications in \secref{discuss}. A similar correlation between stellar mass and absolute magnitude has been reported for LBGs at $z\simeq 5$--$6$ by \citet{star09}. \figref{fig:mass_abs} shows that LBGs at $z\simeq 1.5$--$5$ follow similar linear correlation between stellar mass and UV absolute magnitude (within uncertainties) for $M_{\rm uv}$ between --19 and --22.5~mag. It is important to note that \citet{shap05} does not find any correlation between the stellar mass and UV absolute magnitude for star-forming galaxies at $z\simeq 2$ with stellar masses $\gtrsim\,$10$^{10}$~M$_{\odot}$. This could be due to different color-selection technique (BX/BM) used by the \citet{shap05} to select star-forming galaxies, whose physical properties could differ from the dropout selected LBGs at these masses \citep[e.g.,][]{ly11,habe12}. It is also possible that their sample --- which consists of spectroscopically confirmed bright galaxies with stellar masses greater than or equal to 10$^{10}$~M$_{\odot}$ --- has more massive galaxies than our sample and it is uncertain how massive galaxies would follow this correlation. The ERS observations are too limited in area and depth to cover a larger luminosity range, so we cannot predict how this relation will evolve for luminous ($M_{\rm uv}\! <\,$--22.5~mag) or dwarf ($M_{\rm uv}\!>\,$--19~mag) galaxies at these redshifts. \subsection{SFR vs Stellar Mass}\label{sfr_mass} The correlation between the current SFR and stellar mass in star-forming galaxies, also known as `main sequence of star-formation' (MS), has been observed at $z\lesssim 2$ \citep[e.g.,][]{noes07,elba07,dadd07}. These studies have shown that the MS relation seems to be not evolving strongly with redshift, but the zeropoint does: that is high redshift ($z\simeq 2$) star-forming galaxies are forming stars at a higher rate than similar mass local galaxies. In \figref{fig:mass_sfr}, we investigate this relation and star-formation histories for LBGs at $z\simeq 1.5$--$5$. These quantities are based on best-fit SEDs, and their typical uncertainties are shown in the lower-right corner. The dotted lines are best-fit line obtained by keeping the logarithmic slope fixed at 0.90, estimated for star-forming galaxies at $z\le 2$ \citep[e.g.,][]{elba07,dadd07,sawi12}. The dot-dash lines show the scatter from the best-fit line, which is $\sim$0.6 dex for LBGs at \uvdrops\ and about 0.4 dex for LBGs at \bvdrops. We also obtained the best-fit logarithmic slope for this relation, and found the fitted slope in the range of 0.81$\pm$0.30 for our LBG samples, which is consistent with 0.90 within the estimated 1$\sigma$ scatter in this relation. We find that a proportionality with a logarithmic slope of 0.90 provides a good fit to the data with few outliers at stellar mass greater than 10$^{10}$~M$_{\odot}$. \citet{finl06} have shown that tight relation exists between SFR and stellar mass for galaxies at $z\simeq 4$ using the cosmological hydrodynamic simulations, which is also consistent with the observations \citep{bouw12}. \citet{finl06} also point out that the scatter in the \figref{fig:mass_sfr} could be a measure of SFR `burstiness' as a function of stellar mass. This means that the linear relation (with a logarithmic slope of $\sim$0.90) indicate an average SFR for a given stellar mass, but galaxies can also experience bursts of up to two times the average SFR value at the same stellar mass as shown by the scatter. The scatter in the SFR versus stellar mass relation for LBGs at \uvdrops\ is slightly larger than $\sim$0.3 dex --- observed at $z\simeq 2$ by \citet{dadd07} --- possibly because of few galaxies forming a sharp edge towards high SFR values, as seen in the relation for \wfcfuv- and \wfcnuv-dropouts (upper panel in \figref{fig:mass_sfr}). These galaxies have low stellar ages (less than 10~Myr), which could be highly uncertain as discussed in \secref{age_mass_sfr}. It is also possible that this edge could be an artifact due to lower limits on the model parameters $\tau$ and \emph{t} \citep[e.g.,][]{hain12}. We also note that \citet{mclu11} argue that the tightness in the SFR-stellar mass relation depends on the assumed SFH. The scatter in this relation is much less for a constant SFH, while it is much larger for other SFHs. Therefore, it is also likely that the larger scatter we see in \figref{fig:mass_sfr} could be due to different SFHs. \figref{fig:mass_sfr} shows that, though our data has little more scatter compared to the MS relation at $z\lesssim 2$, the majority of our galaxies fall on to this relation characterized by a logarithmic slope of 0.90. A similar correlation is observed at $z\simeq 6$--$8$ by \citet{mclu11}, and supported by cosmological hydrodynamic simulations of \citet{finl11}. Our observations confirm this MS relation for star-forming galaxies from $z\simeq 1.5$ to $5$, implying that --- on average --- their star-formation histories are similar. \subsection{Implications}\label{discuss} In previous sections, we have shown that LBGs at \uvdrops\ --- on average --- are massive, dustier, and have higher star-formation rates than LBGs at \bvdrops\ with similar luminosities, though it should also be noted that they are not very different within estimated 1$\sigma$ uncertainties. As pointed out by \citet{papo11}, the number densities of galaxies at fixed luminosity could change substantially over this redshift range, which could lead to potential biases when comparing galaxies at different redshifts. However, the general trends we observe in stellar masses, SFRs, and dust extinction are supported by other independent means. The characteristics UV luminosity ($L_{\rm uv}^*$) is increasing as a function of redshift from $z\sim 8$ to $2$ (e.g., H10), which implies increase in SFRs with time, while \citet{fink10} have shown that stellar masses for $M_{\rm uv}^*$ LBGs grow from $z\simeq 8$ to $2$. The UV spectral slope $\beta$ shows evolution as a function of redshift (\figref{fig:beta_z}), which could indicate lower dust content at higher redshifts. The higher dust content in LBGs at lower redshift is also in accordance with the studies at $z\simeq 1$ \citep[e.g.,][]{burg07,basu11}, while the \citet{verm07} supports the lower dust content in LBGs at $z\simeq 5$. Therefore, the ensemble properties of LBGs in our sample are in general agreement with the expected results. The stellar mass--UV luminosity relation (\figref{fig:mass_abs}) and the SFR--stellar mass relation (\figref{fig:mass_sfr}) are based on measurements from best-fit SEDs, therefore, it is possible that these quantities are not totally independent, which might affect their observed correlations. To investigate this, we show distributions of mass-to-light (M/L; Mass/$L_{\rm uv}$) ratios and specific SFRs (SSFR; SFR/Mass) in \figref{fig:ssfr}. The black (red) histograms show distribution for LBGs at \uvdrops\ (\bvdrops), and the median values are shown by dashed vertical lines. The median values of M/L ratio and SSFR for LBGs at \bvdrops\ are slightly lower than that at \uvdrops, but are still consistent within the 1$\sigma$ uncertainties as shown by the error bar on the top of the black histogram. A two-sided K-S test --- in each panel --- indicates a probability \emph{less} than 0.05 that the distributions (red and black histograms) are drawn from the \emph{same} parent distribution. The constancy of the M/L ratio and SSFR between $z\simeq 1.5$ and $5$ agrees very well with the constant slope we find in \figref{fig:mass_abs} and \figref{fig:mass_sfr} for our sample of LBGs, though with a slightly larger scatter. Stellar masses of LBGs at \uvdrops\ are generally well correlated with UV absolute magnitude and current SFR, as expected for star-forming galaxies at similar redshifts \citep[e.g.,][]{elba07,dadd07,sawi12}. These correlations implies very similar mass assembly and SFH for these galaxies, but the exact nature of SFHs is still not clearly understood. \citet{papo11} showed that the cosmologically averaged SFRs of star-forming galaxies at $3 < z < 8$ --- at constant co-moving number density --- increase smoothly from $z=8$ to $3$, and the stellar mass growth in these galaxies is consistent with this derived SFH. The scenario of rising SFH \citep[see also][]{lee10} is also supported by recent results from the cosmological hydrodynamic simulations \citep[e.g.,][]{finl11}. The models with rising SFHs conflicts with the assumptions that the SFR in distant galaxies is either constant or decreasing exponentially with time \citep[e.g.,][]{papo05,shap05,labb10}. Though, we remind the reader that the models with rising SFHs advocated by \citet{papo11} and others correspond to a cosmologically averaged SFHs for typical galaxies, and not individual galaxies, because they could involve random events that changes their instantaneous SFR. \citet{papo11} and \citet{lee10} also argue that rising SFHs are most beneficial to higher redshift ($z\gtrsim 3$) galaxies. We find that for our assumed SED model parameters, the LBGs between redshift $z\simeq 1.5$ and $5$ --- on average --- have similar SFHs, though the precise nature of SFHs at all redshift is still under debate, and could also affect the SFR--stellar mass correlation. Our analysis demonstrates that the dropout selected galaxies at \uvdrops\ --- within luminosities probed here --- show similar correlations between physical parameters (SFR, stellar mass, UV luminosity) as other star-forming galaxies selected using different color criteria (e.g., $sBzK$, BX/BM) at $z\simeq 2$. This is consistent with the \citet{ly11} conclusion that majority ($\sim$80--90\%) of the dropout selected galaxies overlap with other color selected star-forming galaxies with stellar masses less than 10$^{10}$~M$_{\odot}$. The stellar mass range for our current sample is between $\sim$10$^{8}$ and $\sim$10$^{10}$ M$_{\odot}$, with sample completeness around 10$^{9-9.5}$ M$_{\odot}$. Significant differences between the dropout selected sample and other color selected samples of star-forming galaxies at $z\simeq 2$ exists for massive galaxies ($\gtrsim\,$10$^{10}$ M$_{\odot}$; \citealt{ly11}). Therefore, it is vital to use uniform selection technique at all redshifts to avoid any selection biases. The Lyman break dropout technique is the most convenient and widely used method to select galaxies at $z\gtrsim 3$, and we have shown that LBGs at \uvdrops\ selected using this dropout technique have similar physical properties (within uncertainties) as LBGs at \bvdrops\ with similar luminosities. Hence, LBG selection at $z<3$ is important to understand properties of LBGs and properly investigate their evolution as a function of redshift. The validity of LBG properties over wide luminosity and mass range can be investigated in detail with the upcoming and future WFC3 UV surveys such as CANDELS \citep{grog11,koek11} and the WFC3 UV UDF \citep[PI: H. Teplitz]{rafe12}. | 12 | 6 | 1206.6116 |
1206 | 1206.4055_arXiv.txt | We present a detailed analysis (including redshift tomography) of the cosmic dipoles in the Keck+VLT quasar absorber and in the Union2 SnIa samples. We show that the fine structure constant cosmic dipole obtained through the Keck+VLT quasar absorber sample at $4.1\sigma$ level is anomalously aligned with the corresponding dark energy dipole obtained through the Union2 sample at $2\sigma$ level. The angular separation between the two dipole directions is $11.3^\circ \pm 11.8^\circ$. We use Monte Carlo simulations to find the probability of obtaining the observed dipole magnitudes with the observed alignment, in the context of an isotropic cosmological model with no correlation between dark energy and fine structure constant $\alpha$. We find that this probability is less than one part in $10^6$. We propose a simple physical model (extended topological quintessence) which naturally predicts a spherical inhomogeneous distribution for both dark energy density and fine structure constant values. The model is based on the existence of a recently formed giant global monopole with Hubble scale core which also couples non-minimally to electromagnetism. Aligned dipole anisotropies would naturally emerge for an off-centre observer for both the fine structure constant and for dark energy density. This model smoothly reduces to \lcdm for proper limits of its parameters. Two predictions of this model are (a) a correlation between the existence of strong cosmic electromagnetic fields and the value of $\alpha$ and (b) the existence of a dark flow on Hubble scales due to the repulsive gravity of the global defect core (`Great Repulser') aligned with the dark energy and $\alpha$ dipoles. The direction of the dark flow is predicted to be towards the spatial region of lower accelerating expansion. Existing data about the dark flow are consistent with this prediction. | According to the cosmological principle, the Universe is homogeneous and isotropic on scales larger than a few hundred Mpc. The main source of evidence which supports this assumption comes from the Cosmic Microwave Background (CMB) which appears to be isotropic to a high degree up to a dipole term which is assumed to be due to our motion with respect to the CMB frame. However, there has been some recent observational evidence which could be interpreted as a hint for deviations from large scale statistical isotropy. Such evidence includes alignment of low multipoles in the CMB angular power spectrum \cite{Copi:2010na}, large scale velocity flows \cite{Watkins:2008hf,Kashlinsky:2008ut} and large scale alignment in the QSO optical polarisation data \cite{Hutsemekers:2005iz} (see Ref.~\cite{Ciarcelluti:2012pc} for an interesting related theoretical model). These effects appear to persist on scales of 1 Gpc or larger and could constitute early hints for a deviation from the FLRW metric on large cosmological scales and the existence of a cosmological preferred axis. This possibility is further enhanced by the fact that the anisotropy directions implied by these observations appear to be abnormally close to each other~\cite{Antoniou:2010gw}. The above hints for cosmological anisotropy have motivated searches for deviations from the cosmological principle by considering the angular distribution of luminosity distances of Type Ia supernovae (SnIa) in the redshift range $z\in \left[0.015,1.4\right]$ \cite{Antoniou:2010gw,Colin:2010ds,Cooke:2009ws,Blomqvist:2008ud,Cooray:2008qn, Gupta:2010jp,Schwarz:2007wf,Campanelli:2010zx,Cai:2011xs}. Even though all these studies are consistent with isotropy, in many of them, a mild evidence ($1\sigma-2\sigma$) of anisotropic expansion was found \cite{Antoniou:2010gw,Colin:2010ds,Cooke:2009ws,Cai:2011xs,Schwarz:2007wf} mainly coming from low redshift data, while in others \cite{Blomqvist:2008ud,Gupta:2010jp,Campanelli:2010zx} no evidence of anisotropy was found. The inability of the later studies to pick up any anisotropy is perhaps due to the methods and data used which were not sensitive enough to particular types of anisotropy. Additional hints for such possible deviations from the cosmological principle have recently been obtained by the angular distribution of the fine structure constant $\alpha$ in the redshift range $z\in \left[0.2223,4.1798\right]$ as measured by the quasar absorption line spectra using the many multiplet method \cite{King:2012id}. If in the case of SnIa the dipole anisotropy was mild (about $1-2\sigma$), in the case of the fine structure constant the anisotropy has been found to be significant ($4.1 \sigma$). Some earlier studies had claimed possible variation of the fine structure `constant' with time \cite{Webb:1998cq}. This possibility has led to extensive theoretical modelling in the literature so far \cite{varconst,Chiba:2011bz} with emphasis on the possible connection of this variation with dark energy (quintessence)\cite{alpha-quint}. However, there has been comparatively less interest in the possibility of spatial variation of $\alpha$ (see however \cite{Olive:2012ck,Olive:2010vh} for recent studies) and its connection with dark energy. The anisotropy analysis of Ref.~\cite{Antoniou:2010gw} for the SnIa sample was based on the Union2 dataset \cite{Amanullah:2010vv} which consists of $557$ SnIa. A hemisphere comparison method was used to find the hemisphere pair with maximal anisotropy with respect to \lcdm fits. The maximum anisotropy direction was found to be towards $(l,b)= (309^\circ, 18^\circ)$ but the magnitude of this dark energy anisotropy was found to be consistent with statistical isotropy at the $2\sigma$ level. Similar results were obtained in Ref.~\cite{Colin:2010ds} where a redshift tomography also revealed that most of the contribution to the mild dark energy dipole comes from the low redshift SnIa. The anisotropy analysis of the fine structure constant $\alpha$ \cite{King:2012id} is based on a large sample of quasar absorption-line spectra (295 spectra) obtained using UVES (the Ultraviolet and Visual Echelle Spectrograph) on the VLT (Very Large Telescope) in Chile and also previous observations at the Keck Observatory in Hawaii. An apparent variation of $\alpha$ across the sky was found. It was shown to be well fit by an angular dipole model $\daa=A \cos\theta + B$ where $\theta$ is the angle with respect to a preferred axis and $A,B$ are the dipole magnitude and an isotropic monopole term. The dipole axis was found to point in the direction $(l,b)= (331^\circ, -14^\circ)$ and the dipole amplitude $A$ was found to be $A = (0.97\pm 0.21)\times 10^{-5}$. The statistical significance over an isotropic model was found to be at the $4.1\sigma$ level. The analysis of Ref.~\cite{King:2012id} has received criticism \cite{carroll} based mainly on the fact that its quasar sample combines two datasets (Keck and VLT) with different systematic errors which have a small overlapping subset and cover opposite hemispheres on the sky. The axis connecting these two hemispheres has similar direction with the direction of the obtained dipole. The response of the authors of Ref.~\cite{King:2012id} was based on the fact that in the equatorial region of the dipole, where both the Keck and VLT samples contribute a number of absorbers, there is no evidence for inconsistency between Keck and VLT. The controversy about the possible problems in the analysis of Ref.~\cite{King:2012id} and the angular proximity between the dark energy axis of Ref.~\cite{Antoniou:2010gw} and the $\daa$ axis of Ref.~\cite{King:2012id} constitutes the motivation to analyse both the SnIa and the quasar datasets in a similar and consistent manner. Thus we re-analysed both datasets and fit them to the same dipole+monopole ansatz of the form $A \cos\theta + B$. This type of anisotropy fit is different from the corresponding SnIa fits of previous studies. Our goal is to address the following questions: \begin{enumerate} \item What are the best fit dipoles (magnitudes $A$ and directions in galactic coordinates) for the Union2 and Keck+VLT samples? What is the angle between the two dipole directions? \item How likely is it to obtain these dipole magnitudes in the context of an isotropic underlying model? How likely is it to obtain the observed angle between the dipoles if the two underlying models were isotropic and uncorrelated? We address these questions by producing a large number of Monte Carlo isotropic datasets simulating the Union2 and the Keck+VLT samples under the assumption of isotropic and uncorrelated underlying models. We then compare the obtained probability distributions for the dipole magnitudes and angles with the observed magnitudes and angle. \item How do the answers to the above questions change if we consider three different redshift slices (bins) for each dataset (low, medium and high redshift) with approximately equal number of datapoints in each bin? Is there a particular redshift range where the dark energy and the fine structure dipoles are more prominent and how is this range related with the quality of the data in each bin? \end{enumerate} These questions are addressed in detail in the following sections. In particular, the structure of this paper is the following: in the next section we derive the magnitudes and directions of the best fit dark energy and fine structure dipoles for the full Union2 and Keck+VLT datasets thus addressing the above question 1. We also perform $10^4$ Monte Carlo simulations of the Union2 and Keck+VLT datasets based on an isotropic best fit \lcdm model and on an isotropic best fit monopole model respectively. We then use these simulations to address the above question 2. In section~\ref{sec:z_tomography} we perform a redshift tomography to address question 3 and find the redshift range where the dipoles appear to be more pronounced. In section~\ref{sec:physical_mech} we discuss a physical model that could reproduce the observed dipole alignment. Finally, in section~\ref{sec:conclusions} we conclude, summarise our basic results and discuss future prospects of the present work. | \label{sec:conclusions} We have used the Keck+VLT dataset and the Union2 dataset to show that the value of the fine structure constant and the rate of accelerating expansion are better described by coinciding dipoles than by isotropic cosmological models. The key feature of our analysis is that it applies identical method (fit to dipole+monopole anisotropy) to both the Keck+VLT dataset and the Union2 dataset. This consistency, combined with the apparent dipole nature of the anisotropy, has allowed a consistent comparison of the two dipoles. Using Monte Carlo simulations and covariance matrix error estimates, we find that the probability that these coinciding dipoles are both produced in the context of a cosmological model where fine structure constant and dark energy are isotropic and uncorrelated is less than one part in $10^6$. A redshift tomography analysis dividing the two datasets in three redshift bins revealed that the highest data quality redshift bins correspond to low redshifts for the Union2 sample and high redshift for the Keck+VLT data. The dipole direction for the Keck+VLT data depends weakly on redshift while the Union2 dipole direction depends more strongly on redshift and it is the low redshift (and lowest error) bin that is best aligned with the Keck+VLT dipole. The directional uncertainty is significantly larger for the medium and higher redshift Union2 dipoles. It is therefore important to improve the quality of intermediate and high redshift SnIa data in order to further test the alignment of the dark energy dipole with the fine structure constant dipole. An important issue that we have not addressed in the present paper is the effect of systematic errors of the Keck+VLT sample. This issue has been addressed in detail in Ref.~\cite{King:2012id} where no significant source of systematic errors was identified. The main concern has been the possibility of careless merging of the two datasets (VLT and Keck) which in principle have different systematics and effectively cover opposite hemispheres of the sky which coincide with the direction of the identified Keck+VLT dipole. The concern therefore is that the large identified dipole magnitude originates from a hidden difference in systematic errors between the VLT and Keck samples\cite{carroll}. According to Ref.~\cite{King:2012id} this does not appear to be the case for the following reasons: \begin{itemize} \item the dipole directions at high and low redshifts are in agreement (this is confirmed in our study too as shown in Figs.~\ref{fig:DipoleDirs} and \ref{fig:AngDistWebb}); \item the directions of the dipoles fitted by the VLT and by the Keck samples separately are in agreement; \item the absorbers that are common to both the VLT and the Keck sample provide consistent values for $\alpha$. \end{itemize} Even though the above arguments of Ref.~\cite{King:2012id} are reasonable, a truly convincing analysis would involve observation of the same objects with a different telescope. This has already been done by the Subaru telescope in August 2004 \cite{Chiba:2011bz}. An analysis of these observations could provide a particularly useful independent verification of the fine structure constant dipole. Finally we have proposed a theoretical model that has the potential to predict strong aligned dipoles for the fine structure constant and for dark energy. The model is based on a non-minimal coupling of a topologically non-trivial scalar field to electromagnetism (extended topological quintessence). In such a model, an off-centre observer with respect to the Hubble scale core of a global monopole would naturally observe large aligned dipoles for the fine structure constant and dark energy. In fact it should be possible to reconstruct both the scalar field potential and the non-minimal coupling form using the Keck+VLT and the Union2 samples. A robust prediction of the non-minimally coupled defect model is the weak dependence of the value of $\alpha$ on the existence of local strong magnetic fields as discussed at the end of the previous section. Another interesting prediction is the existence of peculiar velocities in the direction away from the center of the global monopole due to the repulsive effects of antigravity (negative pressure) in the defect center. An off-center observer would experience this Hubble scale flow as a dipole dark flow. Such dipole dark flow has indeed been observed \cite{Kashlinsky:2008ut,Watkins:2008hf} and it is attributed to the existence of a Great Attractor which could be present on $Gpc$ scales (perhaps even at a neighboring universe \cite{MersiniHoughton:2008rq}). In our model such a dark flow could be due to a 'Great Repulser' whose role would be played by the core of the Hubble scale non-minimally coupled defect. The predicted direction of such a flow should be away from the defect core (Great Repulser) in the direction of maximum deceleration ($b=-15.1^\circ \pm 11.5^\circ$, $l=309.4^\circ \pm 18.0^\circ$) (see Table \ref{tab:Union2Dipole}). The direction of the observed dark flow is ($b=8^\circ \pm 6^\circ$, $l=287^\circ \pm 9^\circ$)\cite{Watkins:2008hf} which is consistent within $1\sigma$ with the direction of the dark energy and $\alpha$ dipoles. It also points towards the region of lower acceleration as predicted by our model. A robust prediction of our model with respect to the dark flow is that it should reverse direction at large enough redshifts as we start seeing on the 'other side' of the 'Great Repulser' (defect core). A detailed investigation of the consistency of the above predictions with cosmological observations is an interesting extension of the present analysis. {\bf Numerical Analysis Files:} The data, Mathematica and C++program files used for the numerical analysis files may be downloaded from \url{http://leandros.physics.uoi.gr/defsdipoles}. | 12 | 6 | 1206.4055 |
1206 | 1206.4263_arXiv.txt | In this paper we describe the first data release of the the Visible and Infrared Survey Telescope for Astronomy (VISTA) Deep Extragalactic Observations (VIDEO) survey. VIDEO is a $\sim~12$~degree$^{2}$ survey in the near-infrared $Z$,$Y$,$J$,$H$ and $K_{\rm s}$ bands, specifically designed to enable the evolution of galaxies and large structures to be traced as a function of both epoch and environment from the present day out to z=4, and active galactic nuclei (AGN) and the most massive galaxies up to and into the epoch of reionization. With its depth and area, VIDEO will be able to fully explore the period in the Universe where AGN and starburst activity were at their peak and the first galaxy clusters were beginning to virialize. VIDEO therefore offers a unique data set with which to investigate the interplay between AGN, starbursts and environment, and the role of feedback at a time when it was potentially most crucial. We provide data over the VIDEO-XMM3 tile, which also covers the Canada-France-Hawaii-Telescope Legacy Survey Deep-1 field (CFHTLS-D1). The released VIDEO data reach a $5\sigma$ AB-magnitude depth of $Z=25.7$, $Y=24.5$, $J=24.4$, $H=24.1$ and $K_{\rm s}=23.8$ in 2~arcsec diameter apertures (the full depth of $Y=24.6$ will be reached within the full integration time in future releases). The data are compared to previous surveys over this field and we find good astrometric agreement with the Two-Micron All Sky Survey, and source counts in agreement with the recently released UltraVISTA survey data. The addition of the VIDEO data to the CFHTLS-D1 optical data increases the accuracy of photometric redshifts and significantly reduces the fraction of catastrophic outliers over the redshift range $0<z<1$ from 5.8 to 3.1 per cent in the absence of an $i-$band luminosity prior. However, we expect the main improvement in photometric redshifts will come in the redshift range $1<z<4$ due to the sensitivity to the Balmer and 4000\AA\, breaks provided by the near-infrared VISTA filters. All images and catalogues presented in this paper are publicly available through ESO's “phase 3” archive and the VISTA Science Archive. | \label{sec:intro} We are already at a point where we have excellent constraints on the spatial distribution and properties of galaxies in the local Universe from the largest surveys ever undertaken, namely the 2-degree field Galaxy Redshift Survey \citep[2dFGRS; ][]{Colless2001}, and the Sloan Digital Sky Survey \citep[SDSS; ][]{Adelman-McCarthy2008} in the optical, the Two-Micron All Sky Survey \citep[2MASS; ][]{2mass} in the near-infrared, and the IRAS redshift survey \citep{Saunders2000} selected in the far-infrared. Over the past few years these surveys to study the formation and evolution of galaxies in the Universe, and how they trace the large-scale structure of the Universe, have been complemented by deeper and narrower surveys at all wavelengths, carefully balancing depth and area to ensure that representative volumes of the Universe are observed. Most notably the Galaxy and Mass Assembly \citep[GAMA; ][]{Driver2011} Survey covering $\sim 100$ square degrees is building on the legacy of the SDSS and 2dfGRS spectroscopic surveys to obtain accurate redshift information down to an $r-$band limit of $r \sim 19.8$ (corresponding to $z$~$\ltsim\,0.5$), with the UKIRT Infrared Deep Sky Survey Large-Area Survey \citep[UKIDSS-LAS; ][]{Lawrence2007} providing important near-infrared data. At the other extreme, the deepest ground-based surveys such as COSMOS \citep{Scoville2007} and the UKIDSS Ultra Deep Survey coupled with the Subaru-XMM Deep Survey \citep[e.g.][]{Foucaud2007, Furusawa2008}, cover $\sim 1$~degree$^{2}$ fields, relying on photometric redshifts to trace the evolutionary history of galaxies \citep[e.g.][]{Ilbert2009}, coupled with extremely time intensive spectroscopic follow-up\citep[e.g.][]{Lilly2007, lefevre2005}. The Visible and Infrared Survey Telescope for Astronomy \citep[VISTA;][]{Emerson2004} is now taking the surveys carried out within UKIDSS one step further, with its larger field of view and improved sensitivity at the bluer ($Z, Y, J$) wavelengths of the near-infrared window. VISTA is undertaking several near-infrared public surveys, one of which is the VISTA Deep Extragalactic Observations (VIDEO) Survey. The aim of the VIDEO survey is to gain a representative view of the Universe at $0< z < 4$, allowing galaxy evolution to be traced over the majority of the history of the Universe, from the richest clusters to the field. The VIDEO survey fields are chosen to incorporate current and future multi-wavelength data sets to facilitate the broadest exploitation of the new near-infrared survey data both within ESO and across the globe. The main scientific aims of VIDEO are described below. Throughout this paper we use the AB-magnitude system \citep{OkeGunn1983} and adopt the concordance cosmology $\Omega_{M} = 0.3$, $\Omega_{\Lambda} = 0.7$ and $H_{\circ} = 70$\,km\,s$^{-1}$\,Mpc$^{-1}$. | In this paper we have given an overview of the scientific goals, the data processing and the technical details of the VISTA Deep Extragalactic Observations (VIDEO) survey. Our stacked images from our first major data release in the VIDEO-XMM3 field reach $5\sigma$ depths of $Z=25.7$, $Y=24.5$, $J=24.4$, $H=24.1$ and $K_{\rm s}=23.8$ in the AB-mag system in a 2~arcsec diameter aperture, with the expectation that the full depth in the $Y$ band will be reached in the planned integration times. This quality of data will eventually cover 12~degree$^2$ of the best-studied fields in the southern hemisphere, making it a unique survey for investigating the formation and evolution of the rarest and most massive galaxies down to the sub-$L^{\star}$ galaxy population as a function of galaxy environment, over 90 per cent of cosmic time. For this first major data release, we have emphasized the power of combining the VIDEO survey with the CFHTLS-D1 optical data in the VIDEO-XMM3 field for improving photometric redshifts (these can be obtained from the VIDEO consortium web page\footnote{http://star.herts.ac.uk/$\sim$dgb/video}). The addition of the VIDEO bands to optical data results in a considerable decrease in the fraction of catastrophic failures in the photometric redshifts, from 7.4 per cent to 3.8 per cent without an $i-$band luminosity function prior and from 5.4 to 3.3 per cent when the luminosity function prior is included. Furthermore, the accuracy of the photometric redshifts with the inclusion of the near-infrared VIDEO data should increase substantially at $z>1$ due to the ability to bracket strong spectral features, such as the Balmer and 4000\AA\, breaks, over a broader range of redshifts. However, we also emphasize the importance of additional deep optical imaging data over the full VIDEO Survey area. Thus, future observations as part of the Dark Energy Survey and the VLT Survey Telescope (VST) will play a crucial role in the full scientific exploitation of the VIDEO data. The images and catalogues described here will be publicly available from the ESO archive\footnote{http://www.eso.org/sci/observing/phase3/data\_releases/} and the VISTA Science Archive\footnote{http://surveys.roe.ac.uk/vsa}. Access to value-added information, and updates on the progress of the VIDEO survey can also be found at the VIDEO survey website (http://star.herts.ac.uk/$\sim$dgb/video). The forthcoming release at the end of June 2012 will also include $Z$ and $K_{\rm s}-$band data in the ELAIS-S1 northern field, full details of which will accompany the data within the ESO archive and the VSA. The remainder of the VIDEO Survey data and value-added catalogues will be released on regular 6-12 month timescales. A description of the data products will accompany all releases and be held at the ESO and VSA archives, as well as on the VIDEO survey website. | 12 | 6 | 1206.4263 |
1206 | 1206.5008_arXiv.txt | We present SMA and CARMA continuum and spectral line observations of five dense cores located in the Perseus and Ophiuchus molecular clouds whose masses exceed their thermal Jeans masses. Three of these cores have previously been identified as being starless and two have been classified as being possibly protostellar. We find that one core is certainly protostellar. The other four cores, however, are starless and undetected in both \CEO\ and 1.3\,mm continuum emission. These four starless cores have flat density profiles out to at least $\sim$0.006\,pc, which is typical for starless cores in general. Density profiles predicted by some collapse models, especially in the early stages of infall, are consistent with our observations. Archival data reveal that these starless cores have significant non-thermal support against collapse, although they may still be unstable. | Recent interferometric observations of supposedly ``starless'' cores in nearby molecular clouds have found several hidden protostars \citep[e.g.,][]{Pineda11, Dunham11, Schnee12, Chen12}. These observations, however, almost never reveal substructures or evidence for fragmentation in starless cores \citep{Olmi05, Schnee10}, with some notable exceptions, such as R CrA SMM 1A \citep{Chen10} and L183 \citep{Kirk09}. The lack of observed substructure in starless cores may imply that stellar multiplicity begins during the protostellar stage. For instance, in the disk fragmentation theory of multiple star formation, a massive accretion disk around a protostar can become unstable and fragment, creating a binary or higher-order system \citep{Adams89, Bonnell94}. Alternatively, turbulence in starless cores may be forming the seeds of multiplicity \citep{Fisher04, Goodwin04, Goodwin07}, but these seeds may be below the threshold of detectability, as recently suggested by \citet{Offner12}. In this paper, we present interferometric observations of five cores whose masses derived from dust emission at 850\,\micron\ exceed their respective thermal Jeans masses by at least a factor of 2. As such, they are likely to be in the process of forming protostars and/or fragmenting. These cores, given the label ``super-Jeans'', were identified by \citet{Sadavoy10b} and thus make promising targets in the search for substructure in starless cores. | \label{SUMMARY} In this paper, we presented new SMA and CARMA observations of five super-Jeans cores that have been previously classified as being candidate ``starless'' or ``undetermined'' by \citet{Sadavoy10b}. We find, in agreement with previous observations, that the core Per-8 is actually protostellar and is the origin of a collimated bipolar outflow. The cores Oph-1, Oph-2, Per-2, and Per-6 are truly starless, however, with non-detections in the 1.3\,mm continuum corresponding to 3\,$\sigma$ upper limits to the mass of any embedded point source of $\le$0.02\,\solmass. Our simulations suggest that these four cores have flat density profiles out to radii of at least $\sim$0.006--0.01\,pc, in agreement with previous single-dish and interferometric observations of starless cores. Although Oph-1, Oph2, Per-2, and Per-6 are super-Jeans when considering only their thermal support, when non-thermal support is considered their stability against collapse is much more uncertain. | 12 | 6 | 1206.5008 |
1206 | 1206.3304_arXiv.txt | The cosmological origin of both dark and baryonic matter can be explained through a unified mechanism called hylogenesis where baryon and antibaryon number are divided between the visible sector and a GeV-scale hidden sector, while the Universe remains net baryon symmetric. The ``missing'' antibaryons, in the form of exotic hidden states, are the dark matter. We study model-building, cosmological, and phenomenological aspects of this scenario within the framework of supersymmetry, which naturally stabilizes the light hidden sector and electroweak mass scales. Inelastic dark matter scattering on visible matter destroys nucleons, and nucleon decay searches offer a novel avenue for the direct detection of the hidden antibaryonic dark matter sea. | } The cosmological origin of both baryonic matter~\cite{Riotto:1999yt} and dark matter~\cite{Jungman:1995df} remain an important mystery in our understanding of the early Universe. An array of astrophysical observations indicates that a fraction $\Omega_{\rm b} \approx 4.6\%$ of the energy content of the Universe is baryonic matter, while a fraction $\Omega_{\rm DM} \approx21\%$ is dark matter (DM)~\cite{Komatsu:2010fb}. The Standard Model (SM) is incapable of explaining these observations, providing no viable DM candidate, nor a successful mechanism for generating the baryon asymmetry. Cosmology therefore requires new fundamental physics beyond the SM, and it is important to find ways to detect such new physics experimentally. The apparent coincidence between the densities of dark and baryonic matter, given by $\Omega_{\rm DM}/\Omega_{\rm b} \approx 5$, may be a clue that both originated through a unified mechanism. A wide variety of models have been proposed along these lines within the framework of asymmetric DM~\cite{Nussinov:1985xr,Hooper:2004dc,Kitano:2004sv,Agashe:2004bm, Farrar:2005zd,Kohri:2009yn,Shelton:2010ta,Hut:1979xw,Dodelson:1989cq,Kuzmin:1996he,Gu:2007cw,An:2009vq, Davoudiasl:2010am}; see Ref.~\cite{Davoudiasl:2012uw} for a review. In these scenarios, DM carries a conserved global charge, and its relic abundance is determined by its initial chemical potential. Moreover, if the DM charge is related to baryon number ($B$), then the cosmic matter coincidence is naturally explained for $\mathcal{O}(5 \; {\rm GeV})$ DM mass. In this work, we explore model-building, cosmological, and phenomenological aspects of hylogenesis (``matter-genesis''), a unified mechanism for generating dark and baryonic matter simultaneously~\cite{Davoudiasl:2010am,Davoudiasl:2011fj}. Hylogenesis requires new hidden sector states that are neutral under SM gauge interactions but carry non-zero $B$. $\cp$-violating\footnote{$C$ is charge conjugation and $P$ is parity.} out-of-equilibrium decays in the early Universe generate a net $B$ asymmetry among the SM quarks and an equal-and-opposite $B$ asymmetry among the new hidden states. The Universe has zero total baryon number, but for appropriate interaction strengths and particle masses, the respective $B$ charges in the two sectors will never equilibrate, providing an explanation for the observed asymmetry of (visible) baryons. The stable exotic particles carrying the compensating hidden antibaryon number produce the correct abundance of dark matter. Put another way, DM consists of the missing antibaryons. The minimal hylogenesis scenario, described in Refs.~\cite{Davoudiasl:2010am,Davoudiasl:2011fj}, has the following three ingredients: \begin{enumerate} \item DM consists of two states, a complex scalar $\Phi$ and Dirac fermion $\Psi$, each carrying $B=-1/2$. \item A Dirac fermion $X$, carrying $B=1$, that transfers $B$ between quarks and DM through the gauge invariant operators~\cite{Dimopoulos:1987rk} \beq X\,u^c_{Ri}\,d^c_{Rj}\,d^c_{Rk} \, , \quad X \Psi \Phi \label{eq:hylomin} \eeq where $i,j,k$ label generation (color indices and spinor contractions are suppressed). \item An additional $U(1)^\prime$ gauge symmetry that is kinetically mixed with hypercharge and spontaneously broken near the GeV scale, producing a massive $Z^\prime$. \end{enumerate} With these ingredients, hylogenesis proceeds in three stages, which we illustrate schematically in Fig.~\ref{fig:hylopic}: \begin{enumerate} \item Equal ($\cp$-symmetric) densities of $X$ and $\bar{X}$ are created non-thermally, {\it e.g.}, at the end of a moduli-dominated epoch when the Universe is reheated through moduli decay to a temperature $\TRH$ in the range of $5\;\mev\lesssim \TRH \lesssim 100\;\gev \ll m_X$~\cite{Moroi:1999zb}. \item The interactions of Eq.~\eqref{eq:hylomin} allow $X$ to decay to $u_{Ri}\,d_{R_j}\,d_{Rk}$ or $\bar{\Psi}\Phi^*$, and similarly for $\bar{X}$. With at least two flavours of $X$, these decays can violate $\cp$ leading to slightly different partial widths for $X$ relative to $\bar{X}$, and equal-and-opposite asymmetries for visible and hidden baryons. \item Assuming $\Phi$ and $\Psi$ are charged under $U(1)^\prime$, the symmetric densities of hidden particles annihilate away almost completely, with $\Psi\bar{\Psi} \to Z'Z'$ and $\Phi\Phi^* \to Z'Z'$ occurring very efficiently in the hidden sector, followed by $Z'$ decaying to SM states via kinetic mixing. The residual antibaryonic asymmetry of $\Phi$ and $\Psi$ is asymmetric DM. Likewise, the symmetric density of visible baryons and antibaryons annihilates efficiently into SM radiation. \end{enumerate} Both $\Psi$ and $\Phi$ are stable provided $|m_{\Psi}- m_{\Phi}| < m_p+m_e$, and they account for the observed DM density for $m_\Psi + m_\Phi \approx 5 m_p$, implying an allowed mass range $1.7 \lesssim m_{\Psi,\,\Phi} \lesssim 2.9$ GeV. On the phenomenological side, hylogenesis models possess a unique experimental signature: induced nucleon decay (IND), where antibaryonic DM particles scatter inelastically on visible baryons, destroying them and producing energetic mesons. If $X$ couples through the ``neutron portal'' $u_R d_R d_R$, IND produces $\pi$ and $\eta$ final states, while if $X$ couples through the ``hyperon portal'' $u_R d_R s_R$, IND produces $K$ final states. These signatures mimic regular nucleon decay, with effective nucleon lifetimes comparable to or shorter than existing limits; however, present nucleon decay constraints do not apply in general due to the different final state kinematics of IND. Searching for IND in nucleon decay searches, such as the Super-Kamiokande experiment~\cite{Kobayashi:2005pe} and future experiments~\cite{Akiri:2011dv,IceCube:2011ac,Hewett:2012ns}, therefore offers a novel and unexplored means for discovering DM. \begin{figure}[ttt] \begin{center} \vspace{1cm} \if\withFigures1 \includegraphics[width = 0.18\textwidth]{xbg-step1-eps-converted-to.pdf}~~~~~ \includegraphics[width = 0.36\textwidth]{xbg-step2-eps-converted-to.pdf}~~~~~ \includegraphics[width = 0.36\textwidth]{xbg-step3-eps-converted-to.pdf} \fi \end{center} \caption{(color online) The three steps of hylogenesis.} \label{fig:hylopic} \end{figure} Although the minimal hylogenesis model described above successfully generates the cosmological baryon and DM densities, two puzzles remain. Is there a natural framework to consider DM as a quasi-degenerate scalar/fermion pair? Is there a mechanism to ensure the quantum stability of the GeV-scale masses for hidden sector scalars? Supersymmetry~(SUSY) can provide answers to both questions; the DM pair $(\Phi,\Psi)$ forms a supermultiplet with $B=-1/2$, and the stability of the GeV-scale hidden sector and the $(\Phi,\Psi)$ mass splitting is ensured naturally, provided SUSY breaking is suppressed in the hidden sector compared to the visible sector. The goal of this work is to embed hylogenesis in a supersymmetric framework of natural electroweak and hidden symmetry breaking, and to study in detail the cosmological and phenomenological consequences. In Section~\ref{sec:susy}, we present a minimal supersymmetric extension of the hylogenesis theory described above. We also address the origin of the nonrenormalizable nucleon portal operator $X\,u^c_{Ri}\,d^c_{Rj}\,d^c_{Rk}$. In Section~\ref{sec:baryo}, we investigate the cosmological dynamics of supersymmetric hylogenesis, showing explicitly the range of masses and parameters that can explain the correct matter densities. Section~\ref{sec:susybusting} contains a discussion of how such parameter values can arise in a natural way from various mechanisms for supersymmetry breaking. In Section~\ref{sec:pheno} we investigate the phenomenology of our model, including IND signatures, collider probes, and DM direct detection. Our results are summarized in Section~\ref{sec:conc}. In Appendix~\ref{sec:hportal}, we also present an alternative supersymmetric model based on Higgs portal mixing. | } Through the mechanism of hylogenesis, the cosmological densities of visible and dark matter may share a unified origin. Out-of-equilibrium decays during a low-temperature reheating epoch generate the visible baryon asymmetry, and an equal antibaryon asymmetry among GeV-scale hidden sector states. The hidden antibaryons are weakly coupled to the SM and are the dark matter in the Universe. We have embedded hylogenesis in a supersymmetric framework. By virtue of its weak couplings to the SM, SUSY-breaking is sequestered from the hidden sector, thereby stabilizing its GeV mass scale. The DM consists of two states, a quasi-degenerate scalar-fermion pair of superpartners. We studied in detail one particular realization of supersymmetric hylogenesis, considering several aspects: \begin{itemize} \item We constructed a minimal supersymmetric model for hylogenesis. Hidden sector baryons are chiral superfields $X$ and $Y$, with $B=1$ and $-1/2$, respectively. The lightest $Y$ states are DM, while $X$ decays in the early Universe generate the $B$ asymmetries. \item In addition, we introduced a vector-like $SU(3)_C$ triplet to mediate $B$ transfer between visible and hidden sectors, and a hidden $Z^\prime$ gauge boson (with kinetic mixing) to deplete efficiently the symmetric DM densities. \item We showed that hylogenesis can successfully generate the observed $B$ asymmetry during reheating. We computed the $\cp$ asymmetry from $X$ decay and solved the coupled Boltzmann equations describing the cosmological dynamics of hylogenesis. \item We studied how SUSY breaking is communicated between the visible and hidden sectors through RG effects. We also examined predictions within an AMSB framework. While anomaly mediation explains the late-time reheating epoch from moduli decay, we have not explicitly addressed the issues of tachyonic slepton masses in the visible sector and the origin of SUSY mass terms in the hidden sector. \item Antibaryonic DM annihilates visible nucleons, causing induced nucleon decay to kaon final states, with effective nucleon lifetime in the range $10^{30} - 10^{36}$ years. DM can be discovered in current nucleon decay searches, and this signal remains unexplored. \item Collider searches for monojets and dijet resonances provide the strongest direct constraints on our model, and these signals are correlated with IND. Lifetimes of $10^{30} - 10^{36}$ years correspond to energy scales $\Lambda_{\rm IND} \sim 0.5 - 5$ TeV that can be probed at the LHC. \item DM direct detection experiments and precision searches for hidden photons constrain the $Z'$ kinetic mixing, although our model remains consistent with current bounds. \end{itemize} We emphasize that our specific model was constructed to illustrate general features of hylogenesis, and certainly there are many other model-building possibilities along these lines. Nevertheless, it is clear that supersymmetric hylogenesis provides a technically natural and viable scenario for the genesis of matter, explaining the cosmic coincidence between the dark and visible matter densities and predicting new experimental signatures to be explored in colliders and nucleon decay searches. | 12 | 6 | 1206.3304 |
1206 | 1206.4524_arXiv.txt | The probability density function (PDF) of the gas density in subsonic and supersonic, isothermal, driven turbulence is analysed with a systematic set of hydrodynamical grid simulations with resolutions up to $1024^3$ cells. We performed a series of numerical experiments with root mean square (r.m.s.) Mach number $\mathcal{M}$ ranging from the nearly incompressible, subsonic ($\mathcal{M}=0.1$) to the highly compressible, supersonic ($\mathcal{M}=15$) regime. We study the influence of two extreme cases for the driving mechanism by applying a purely solenoidal (divergence-free) and a purely compressive (curl-free) forcing field to drive the turbulence. We find that our measurements fit the linear relation between the r.m.s.~Mach number and the standard deviation of the density distribution in a wide range of Mach numbers, where the proportionality constant depends on the type of the forcing. In addition, we propose a new linear relation between the standard deviation of the density distribution $\sigma_{\rho}$ and the standard deviation of the velocity in compressible modes, i.e.~the compressible component of the r.m.s.~Mach number $\mathcal{M}_{\mathrm{comp}}$. In this relation the influence of the forcing is significantly reduced, suggesting a linear relation between $\sigma_{\rho}$ and $\mathcal{M}_{\mathrm{comp}}$, independent of the forcing, ranging from the subsonic to the supersonic regime. | Understanding the intricate interplay between interstellar turbulence and self-gravity is one of the key problems in star formation theory. The supersonic turbulent velocity field is likely responsible for the complex and filamentary density structures observed in molecular clouds. It creates dense regions that can become gravitationally unstable and collapse into dense cores and eventually turn into new stars \citep{Elmegreen2004, MacLow2004, McKee2007}. Statistical quantities, describing this process, such as the initial mass function, the core mass function \citep{Padoan2002, Hennebelle2008, Hennebelle2009}, and the star formation rate \citep{Hennebelle2011, Padoan2011} depend on the standard deviation (std.~dev.) of the density of the molecular cloud. The pioneering works by \citet{Padoan1997} and \citet{Passot1998} have shown that the std.~dev.~$\sigma_{\rho}$ of the probability density function(PDF) of the mass-density grows proportional to the root mean square (r.m.s.) Mach number $\mathcal{M}$ of the turbulent flow \begin{equation} \sigma_{\rho}/\! \left<\rho \right>_V = b \,\mathcal{M}\,, \label{eq:Passot1998} \end{equation} where $\left< \rho \right>_V$ is the volume-weighted mean density and $b$ is a proportionality constant. \citet{Federrath2008, Federrath2010} explained the dependence of $\sigma_{\rho}$ on $b$ by taking the modes of the forcing into account that drive the turbulent velocity field. This model predicts for purely solenoidal forcing $b=1/3$ and for purely compressive forcing $b=1$, and explains the large deviations of $b$ ranging from $b=0.26$ to $b=1.05$ in previous works \citep[e.g.][]{Padoan1997, Passot1998, Li2003, Kritsuk2007, Beetz2008, Schmidt2009, Price2011, Konstandin2012, Molina2012}. We follow up on this work and discuss the physical origin of this dependency and introduce a new relation, similar to equation (\ref{eq:Passot1998}), however, correlating the compressible component of the r.m.s~Mach number $\mathcal{M}_{\mathrm{comp}}$ with $\sigma_{\rho}$.\\ In section \ref{sec:Simulation} we explain our numerical setup. We analyse the influence of measuring mass-weighted and volume-weighted distributions in section \ref{subsec:VW_vs_MW}, the influence of the resolution on our measurements in section \ref{subsec:Resolution} and the PDFs of the mass density and the compressible part of the velocity field in section \ref{subsec:The_PDF_of_comressive_modes}. In section \ref{subsec:Relation}, we present the linear relations between the std.~dev.~of the mass density and the r.m.s.~Mach number. In section \ref{subsec:Physical_origin} we discuss the new relation between the std.~dev.~of the mass density and the std.~dev.~of the compressible part of the velocity field. A summary of our results and conclusions is given in section \ref{sec:Summary}. | \label{sec:Summary} We have investigated the influence of solenoidal (divergence-free) and compressive (curl-free) forcing on the PDF of the mass density in subsonic and supersonic turbulence with a set of three-dimensional numerical simulations. We analysed the relation between the std.~dev.~of the mass density distribution and the r.m.s.~Mach number. We found a new relation between the std.~dev.~of the mass density and the std.~dev.~of the compressible part of the velocity field. Our main results are as follows: \begin{itemize} \item Compressive forcing yields mass density PDFs with std.~dev.s proportional to the r.m.s.~Mach number with $b=1$. For solenoidal forcing, we measure $b=1/3$ in the supersonic regime. Our findings are in agreement with previous studies, which however only explored different subsets of the full parameter space investigated here. We also found deviations of our measurements from the linear relation with solenoidal forcing in the subsonic regime. These deviations from the linear relation can be explained with sound waves, which damp the faint compressible velocities and prevent the medium from producing over-densities. \item We found a unique relation between the std.~dev.~of the mass density and the compressible modes of the velocity field with a fit to our data. Our new relation is independent of the driving mechanism and still holds in the subsonic regime, where the flow is mainly influenced by sound waves. It does not show a strong influence on the resolution and other effects, which may cause a non-Gaussian distribution of the density. \item Our relation enables us for the first time to measure the kinetic energy in compressible modes in units of the sound speed, without knowing the r.m.s.~Mach number, the driving mechanism or the sound speed of the medium. This measurement can be used to distinguish between subsonic and supersonic compressive turbulent motions. It will in principle allow us to measure the composition of the kinetic energy in the interstellar medium by combining independent measurements of the total r.m.s.~Mach number \citep[e.g.][]{Burkhart2009} and the std.~dev.~of the density distribution \citep[][]{Brunt2010a, Brunt2010b, Schneider2012}. \end{itemize} | 12 | 6 | 1206.4524 |
1206 | 1206.3132_arXiv.txt | Version 2.0 of CRPropa\footnote{CRPropa is published under the 3rd version of the GNU General Public License (GPLv3). It is available, together with a detailed documentation of the code, at \url{https://crpropa.desy.de}.} is public software to model the extra-galactic propagation of ultra-high energy nuclei of atomic number $Z\leq 26$ through structured magnetic fields and ambient photon backgrounds taking into account all relevant particle interactions. CRPropa covers the energy range $6\times 10^{16} < E/\rm{eV} < A\times 10^{22}$ where $A$ is the nuclear mass number. CRPropa can also be used to track secondary $\gamma$-rays and neutrinos which allows the study of their link with the charged primary nuclei -- the so called multi-messenger connection. After a general introduction we present several sample applications of current interest concerning the physics of extragalactic ultra-high energy radiation. | Cosmic rays are ionized atomic nuclei reaching the Earth from outside the Solar System with energies that exceed $10^{20}~\rm eV$. Although ultra-high energy cosmic rays (UHECRs) were originally discovered in 1939, their sources and propagation mechanisms are still a subject of intense research. During the last decade significant progress has been made due to the advent of high quality and high statistics data from a new generation of large scale observatories. Observables of prime interest are the energy spectrum, mass composition and arrival direction of cosmic rays. A flux-suppression in the energy spectrum above $E \sim 5 \cdot 10^{19}$\,eV has been observed by the HiRes and Pierre Auger Observatories \cite{2008PhRvL.101f1101A,Abbasi-08} and possibly also by the Telescope Array \cite{Matthews:2011zz} indicating either observation of the GZK-effect \cite{Greisen1964,ZatsepinKutzmin1966} or the limiting energy of the sources. Moreover, data of the Pierre Auger Observatory indicate that the arrival directions of the highest energy cosmic rays are correlated with the depth of nearby Active Galactic Nuclei (AGN) or more generally with the nearby extra-galactic matter distribution \cite{Auger2007}. Additionally, measurements of the position of the shower maximum and its fluctuations by the Pierre Auger Collaboration suggest a significant fraction of heavy primaries above $10^{19}$\,eV \cite{Abraham:2010yv}. However, HiRes \cite{2010PhRvL.104p1101A} and preliminary data of the Telescope Array \cite{Matthews:2011zz}), suggest a proton dominance in the same energy range. Unfortunately, the limited number of observed events does not yet allow the extension of these measurements to the aforementioned cutoff energy. Independently of the mass composition, it is not uniquely settled yet if this flux depression is due to energy loss or maximum energy limitation of the sources. Clearly, a better understanding of all of these features and of the effects of cosmic ray propagation through the local Universe is mandatory. UHECRs do not propagate freely in the Inter Galactic Medium (IGM). During their propagation they suffer from catastrophic energy losses in reactions with the intergalactic background light and are deflected by poorly known magnetic fields. Thus, the effects of propagation alter the cosmic ray spectrum and composition injected by sources in the IGM and form the features detected by UHECR observatories. In order to establish the origin of UHECRs, it is of prime interest to quantitatively understand the imprint of the propagation and to disentangle it from the properties of the cosmic rays at their sources. In this respect, it is essential to compare the measured UHECR spectrum, composition and anisotropy with model predictions. This requires extensive simulations of the propagation of UHE nuclei and their secondaries within a given scenario. In particular, the observation that UHECRs may consist of a significant fraction of heavy nuclei challenges UHECR model predictions and propagation simulations. Indeed, compared to the case of ultra-high energy (UHE) nucleons, the propagation of nuclei leads to larger deflections in cosmic magnetic fields and additional particle interactions have to be taken into account, namely, photodisintegration and nuclear decay. To provide the community with a versatile simulation tool we present in this paper a publicly available Monte Carlo code called CRPropa~2.0 which allows one to simulate the propagation of UHE nuclei in realistic one- (1D) and three-dimensional (3D) scenarios taking into account all relevant particle interactions and magnetic deflections. To this end, we extended the former version 1.4 of CRPropa, which was restricted to nucleon primaries, to the propagation of UHE nuclei. CRPropa~1.4 provided an excellent basis for this effort as many of its features could be carried over to the case of UHE nuclei propagation. In the present paper, which accompanies the public release of CRPropa 2.0, the underlying physical and numerical frameworks of the implementation of nuclei propagation are introduced. For technical details the reader is referred to the documentation distributed along with this framework. This paper is organized as follows: Section \ref{sct:CRPRopa1} starts with a short introduction of the publicly available previous CRPropa\,1.4. The extensions which were implemented for nuclei interactions in CRPropa\,2.0 are the subject of section \ref{sct:CRPropa:Interactions}. Section~\ref{Sec:PropAlgo} describes the general propagation algorithm and in section~\ref{sct:Applications} example applications of nuclei propagation with CRPropa are presented. We present a short summary and an outlook in Section~6. Unless stated otherwise, we use natural units $\hbar = c = 1$ throughout this paper. | In the present paper we have presented the new version of our UHECR propagation code CRPropa, a numerical tool to study the effect of extragalactic propagation on the spectrum, chemical composition and distribution of arrival directions of UHECRs on Earth. The main new feature introduced in this new version 2.0 is the propagation of UHE nuclei, taking also into account their interactions with the IGM, in particular photodisintegration, which is modeled according to the numerical framework TALYS. As photodisintegration introduced many more interaction channels than were present in the previous version of CRPropa, we needed to substantially improve the propagation algorithm both in efficiency and in accuracy. We also updated the default model for the extragalactic infrared light to a more recent one. CRPropa\,2.0 can now be used to compute the main observable quantities related to UHECR propagation with the accuracy required by present data: particle spectra, mass composition and arrival direction on Earth, for highly customizable realizations of the IGM, including source distributions and magnetic fields. In addition, the spectra of secondary neutrinos and electromagnetic cascades can be computed down to MeV energies. One of the extensions planned for the future is to include a module to take into account the effect of galactic magnetic fields on deflections. | 12 | 6 | 1206.3132 |
1206 | 1206.3845_arXiv.txt | We have modelled the inner surface brightness profiles of 39 alleged `core' galaxies with the core-S\'ersic model, and provide new physical parameters for the largest ever sample of `core' galaxies fit with this model. When present, additional nuclear components were simultaneously modelled and the typical rms scatter of the fits (out to $\sim$10$\arcsec$) is 0.02 mag arcsec$^{-2}$. Model-independent estimates of each core's break radius are shown to agree with those from the core-S\'ersic model, and a comparison with the Nuker model is provided. We found an absence of cores in what amounts to 18\% of the sample which are reclassified here as S\'ersic galaxies with low values of $n~(\la 4$) and thus shallow inner profile slopes. In general, galaxies with $n<3$ and $\sigma < 183$ km s$^{-1}$ do not have depleted cores. We derive updated relations between core-S\'ersic break radii, their associated surface brightness, bulge luminosity, central velocity dispersion, and predicted black hole mass for galaxies with depleted cores. With the possible exception of NGC 584, we confirm that the inner negative logarithmic profile slopes $\gamma$ are $\la$ 0.3 for the `core' galaxies, and $0 > \gamma > -0.1$ for six of these. Finally, the central stellar mass deficits are found to have values typically within a factor of 4 of the expected central black hole mass. | The stellar distributions in galaxies have played a valuable role in guiding our understanding of the galaxies themselves. In particular, the accessibility of high-resolution imaging offered by the \emph{Hubble Space Telescope} (\emph{HST}) substantially advanced our appreciation of the complexity of galaxy cores (e.g.\ Crane et al.\ 1993; Kormendy et al.\ 1994; Jaffe et al.\ 1994; Ferrarese et al.\ 1994; Grillmair et al.\ 1994; van den Bosch et al.\ 1994; Lauer et al.\ 1995; Byun et al.\ 1996; Gebhardt et al.\ 1996; Carollo et al.\ 1997; Faber et al.\ 1997). For instance, the centers of real galaxies may contain such distinct components as bright active galactic nuclei (AGN), nuclear star clusters, flattened nuclear discs and bars, dust lanes and clouds. On the other hand, giant stellar evacuation zones are also observed. Luminous galaxies with such shallow cores had of course long been known to exist from ground-based observations (e.g.\ King \& Minkowski 1966, 1972; King 1978; Young et al.\ 1978; Binney \& Mamon 1982; see the review by Graham 2012a) but \emph{HST} enabled us to accurately quantify these. After studying 14 bright elliptical galaxies with the pre-refurbished \emph{HST}/WFPC1, Ferrarese et al.\ (1994) introduced a 4-parameter double power law model to describe the inner surface brightness distributions of bright galaxies. While the (relatively brighter) galaxies in their sample which possessed depleted cores with shallow inner profiles were grouped as ``Type I", the remaining galaxies, labeled ``Type II", had a profile that remained steep all the way into the center. Examining a larger sample of galaxies imaged using the same \emph{HST}/WFPC1 high-resolution Planetary Camera, Kormendy et al.\ (1994) and Lauer et al.\ (1995) largely agreed with the division of galaxies presented in Ferrarese et al.\ (1994), referring to them as `core' galaxies and `power-law' galaxies, respectively. They also advanced a double power law model which they dubbed the `Nuker law' for fitting the (underlying host galaxy) surface brightness profiles of early-type galaxies. The Nuker model had an additional fifth parameter to moderate the transition between the two power laws --- as introduced by Hernquist (1990, his eq.\ 43) for modelling the internal density profiles of galaxies. The physical process(es) responsible for the observed difference between the inner surface brightness profiles of `core' galaxies and the fainter `power-law' galaxies (nowadays referred to as `S\'ersic' galaxies because these spheroids have S\'ersic light profiles rather than power-law light profiles) provide valuable clues about the galaxies' past history. In bright galaxies, the widely advocated `dry' (i.e.\ gas poor) galaxy merger hypothesis (e.g. Faber et al.\ 1997) can result in the gravitational sling shot of central stars (core scouring) due to the coalescence of supermassive black holes (SMBHs) from the progenitor galaxies (e.g.\ Begelman, Blandford \& Rees 1980; Makino \& Ebisuzaki 1996; Merritt \& Milosavljevi\'c 2005; Merritt 2006). It is possible that the sizes and mass deficits of such partially depleted cores may reflect the amount of merging and damage caused by the black holes (after having eroded any pre-existing nuclear stellar components: Bekki \& Graham 2010). Having an accurate quantification of the physical parameters defining the centers of galaxies is therefore important. Moreover, reliable break radii $R_b$, used to denote the sizes of the cores, may even be useful for predicting black hole masses (Lauer et al.\ 2007a). While investigating the lack of any connection between the double power-law model and the curved galaxy brightness profiles observed outside of the cores, Graham et al.\ (2003, see their figures~2--4) revealed that the Nuker model's break radius, and other parameters, were not robust quantities but are sensitive to the radial range of the surface brightness profile that is fitted. For example, the break radii $R_b$ were shown to vary by more than a factor of three. The parameters' sensitivity was recognised to arise from the Nuker model's efforts to fit an outer power-law to what is actually a curved brightness profile. The luminosity profiles of bright (core) galaxies ($M_{B} \la -20.5 $ mag) --- which show a downward deviation from the inward extrapolation of their outer S\'ersic (1963, 1968) profile --- were subsequently shown to be precisely represented by the core-S\'ersic model (Graham et al.\ 2003; Trujillo et al. 2004). While Lauer et al.\ (2005) missed this development, Ferrarese et al.\ (2006) found the core-S\'ersic model to be highly applicable to bright early-type galaxies in the Virgo cluster. Lauer et al.\ (2007a,b) subsequently wrote that ``Graham et al.\ (2003) have criticized the Nuker $r_b$ as being sensitive to the domain over which the Nuker law was fitted, particularly when the outer limit of the fit extends only slightly beyond $r_b$. In practice, however, the Nuker laws are fitted over a large radial range that extends well beyond $r_b$.'' However this was not the problem identified by Graham et al.\ (2003), who had demonstrated that the Nuker model parameters deviated further from the true values as the fitted radial extent was {\it increased}. Based on ``work in preparation'' Lauer et al.\ (2007b) refuted that their Nuker break radii were biased ``in any way'' because their radii reportedly agreed very well with model-independent values of where the curvature in the surface brightness profile was a maximum. This was a surprising claim because these latter values should not be dependent on the radial extent of the data while the Nuker model break radii {\it are} a strong function of the fitted radial extent (Graham et al.\ 2003). Kormendy et al.\ (2009, their section 4.1) subsequently buoyed the Nuker model and dismissed the core-S\'ersic model. G\"ultekin et al.\ (2009) then over-looked any and all concerns about the Nuker model which they presented along with Nuker model parameters from Lauer et al.\ (2005), and encouraged readers to use this data, additionally noting that the surface brightness profiles that were fit with the Nuker model are available at the Nuker web page\footnote{http://www.noao.edu/noao/staff/lauer/wfpc2\_profs/}. G\"ultekin et al.\ (2011) continued in this vein, motivating us to further investigate, nearly a decade on, the issue of whether the Nuker model parameters are reliable, physically meaningful quantities, or if instead the core-S\'ersic model parameters may be preferable. At stake is not only the accuracy to which we quantify the cores of galaxies, but our subsequent understanding of cores and how they relate to their galaxy at large. In this paper, we focus on the nuclear structure of galaxies by re-analyzing the surface brightness profiles of all 39 `core' galaxies imaged with the WFC2 / F555W or F606W filter and listed in Lauer et al.\ (2005) to be a `core' galaxy (see section~\ref{Sec_Data}). For reference, Trujillo et al.\ (2004) modelled only 9 possible `core' galaxies, Ferrarese et al.\ (2006) modelled 10, and Richings, Uttley \& Kr$\ddot{\textrm{o}}$ding (2011) have very recently modelled 21 `core' galaxies. We are therefore modelling the largest sample of suspected `core' galaxies to date. For comparison's sake with the Nuker model break radii, we use exactly the same surface brightness profiles as Lauer et al.\ (2005), available at the previously mentioned Nuker web-page. We first concentrate on measuring the core size using the core-S\'ersic model (see sections~\ref{Sec_cS} and \ref{Sec_Fit}). We additionally take the nuclear excess, usually nuclear star clusters or AGN emission, into account while modeling the underlying host galaxy light. In section~\ref{Sec_Rb} we use two model-independent core size estimators and reveal that one of these can not be used while the other is consistent with our core-S\'ersic break radii. Furthermore, we confirm that the published Nuker model break radii are typically 100\% larger than the break in the surface brightness profile determined relative to the inward extrapolation of the outer S\'ersic function (Trujillo et al.\ 2004). We additionally report that `artificial' break radii have been reported in what were alleged to be `core' galaxies but are actually S\'ersic galaxies with no break in their S\'ersic profile and which thus have no partially depleted core relative to their outer light profile (section~\ref{Sec_core}). Throughout this paper we use terms such as 'actual', `true' and `real' break radii and cores when referring to galaxies that have inner surface brightness profiles which break downward from (i.e.\ have lower flux than) the inward extrapolation of the outer S\'ersic model which describes their outer stellar distribution. Sets of structural parameter relations encompassing central as well as global properties are presented in section~\ref{Sec_rel}. In particular, equations involving the break radius and associated surface brightness, and the luminosity, are derived. We also investigate the (core size)-(central black hole mass) relation in section~\ref{Sec_Rb-BH}. Using updated data, we find that the break radius {\it can} be used to consistently predict the black hole mass when using either the $M_{bh}$-$\sigma$ or $M_{bh}$-$L$ relations for `core' galaxies when coupled with our updated $R_{b}$-$\sigma$ and $R_{b}$-$L$ relations. We go on to discuss the detection of additional nuclear components in the full sample of S\'ersic and core-S\'ersic galaxies in section~\ref{Sec_Add} while section~\ref{Sec_Con} summarises our main conclusions. | \label{Sec_Con} We have re-modelled the major-axis, surface brightness profiles of 39 alleged `core' galaxies from Lauer et al.\ (2005), using S\'ersic and core-S\'ersic models. We have additionally and simultaneously accounted for the point sources and additional nuclear components that were excluded by the Nuker analysis. Consistent with earlier published works, we found that the S\'ersic and core-S\'ersic models yield a robust representation of the underlying light distributions of S\'ersic and core-S\'ersic galaxies, respectively, all the way to the resolution limit. The typical rms residual scatter is 0.02 mag arcsec$^{-2}$. The main results of this work are:\\ 1. We have identified 7 of the 39 `core' galaxies from Lauer et al.\ (2005) to be S\'ersic galaxies which do not have partially depleted cores relative to the inward extrapolation of their outer S\'ersic light profile. This situation tends to arise in galaxies and bulges fainter than $M_V \approx -21$ mag. Such galaxies with spheroid S\'ersic index $n ~ \la 3$ or velocity dispersion $\sigma ~ \la 183$ km s$^{-1}$ are not likely to have partially depleted stellar cores. 2. We provide physical parameters ($R_b, \mu_b, \gamma$) for the cores of 32 `core' galaxies , derived using the core-S\'ersic model. 3. Due to noise or real small scale structure, non-parametric core size estimations obtained by locating the maximum of the second logarithmic derivative of the (non-smoothed) light profile, i.e.\ the point of greatest curvature, appear to be unreliable (Fig.~\ref{Fig6}). 4. As with the Nuker model, the break radius of the core-S\'ersic model is shown to coincide with the radius where it has a maximum in the second logarithmic derivative (Fig.\ 9, Right). 5. For the first time, the radius where the negative logarithmic slope of the light profile $\gamma'=1/2$, considered to be a suitable estimator for the size of the core, is shown to be consistent with the core-S\'ersic model break radius (Fig.~\ref{Fig8}). It should, however, be noted that even galaxies without depleted cores will have a radius where $\gamma'$ equals 1/2. Therefore, this measurement cannot be used to identify `true' depleted-core radii. 6. We have compared the core-S\'ersic break radii with the Nuker break radii. In line with previous works, we found that the Nuker break radii are larger than the core-S\'ersic break radii and also $R_{\gamma'=1/2}$: on average, the Nuker break radii are $\sim$2 times bigger than the core-S\'ersic break radii. Furthermore, the surface brightnesses ($\mu_{b}$) at the Nuker model's break radii are up to 2 mag arcsec$^{-2}$ fainter than the surface brightness at the core-S\'ersic model break radii. 7. We have updated various structural parameter relations after excluding galaxies which do not have `real' cores, and using core-S\'ersic parameters. We have also used the bulge magnitude instead of the galaxy magnitude for the disk galaxies. We provide updated $R_{b}$-$L$, $R_{b}$-$\sigma$, $R_{b}$-$\mu_{b}$, $\mu_{b}$-$L$ and $\mu_{b}$-$\sigma$ relations in section 7.1. 8. In contrast to Lauer et al.\ (2007a), we found consistency among three linear $R_{b}$-$M_{BH}$ relationships (section~\ref{Sec_Rb-BH}). While one of these is obtained directly from $R_{b}$ and $M_{BH}$ data (Eq.\ 12), the other two are constructed by combining the $R_{b}$-$\sigma$ and $M_{BH}$-$\sigma$ relations and the $R_{b}$-$L$ and $M_{BH}$-$L$ relations. 9. We detected additional nuclear light in 12 of the 39 sample galaxies. While our sample is rich in `core' galaxies (32/39), 5 of the 12 nucleated galaxies are S\'ersic galaxies: 1 with nonthermal emission from an AGN and 4 with excess stellar light. Five of the 7 nucleated `core' galaxies have AGN emission. These results are in good agreement with previous estimates (e.g.\ Rest et al.\ 2001; C\^ot\'e et al.\ 2006). 10. Following Graham (2004), we derived a tentative central mass deficit for our `core' galaxies using Eq.\ A19 from Trujillo et al.\ (2004). These deficits are about 0.5 to 4 times the expected central supermassive black hole mass. | 12 | 6 | 1206.3845 |
1206 | 1206.1071_arXiv.txt | We report on {\em Chandra} observations of 18 hard X-ray ($>$20\,keV) sources discovered with the {\em INTEGRAL} satellite near the Galactic plane. For 14 of the {\em INTEGRAL} sources, we have uncovered one or two potential {\em Chandra} counterparts per source. These provide soft X-ray (0.3--10\,keV) spectra and sub-arcsecond localizations, which we use to identify counterparts at other wavelengths, providing information about the nature of each source. Despite the fact that all of the sources are within 5$^{\circ}$ of the plane, four of the IGR sources are AGN (IGR~J01545+6437, IGR~J15391--5307, IGR~J15415--5029, and IGR~J21565+5948) and four others are likely AGN (IGR~J03103+5706, IGR~J09189--4418, IGR~J16413--4046, and IGR~J16560--4958) based on each of them having a strong IR excess and/or extended optical or near-IR emission. We compare the X-ray and near-IR fluxes of this group of sources to those of AGN selected by their 2--10\,keV emission in previous studies and find that these IGR AGN are in the range of typical values. There is evidence in favor of four of the sources being Galactic (IGR~J12489--6243, IGR~J15293--5609, IGR~J16173--5023, and IGR~J16206--5253), but only IGR~J15293--5609 is confirmed as a Galactic source as it has a unique {\em Chandra} counterpart and a parallax measurement from previous optical observations that puts its distance at $1.56\pm 0.12$\,kpc. The 0.3--10\,keV luminosity for this source is $(1.4^{+1.0}_{-0.4})\times 10^{32}$\,erg\,s$^{-1}$, and its optical/IR spectral energy distribution is well described by a blackbody with a temperature of 4200--7000\,K and a radius of 12.0--16.4\,\Rsun. These values suggest that IGR~J15293--5609 is a symbiotic binary with an early K-type giant and a white dwarf accretor. We also obtained likely {\em Chandra} identifications for IGR~J13402--6428 and IGR~J15368--5102, but follow-up observations are required to constrain their source types. | The hard X-ray imaging of the Galactic plane by the {\em INTErnational Gamma-Ray Astrophysics Laboratory (INTEGRAL)} satellite continues to find new ``IGR'' sources at energies $>$20\,keV. While {\em INTEGRAL} excels at discovering new sources in the 20--40\,keV band, especially due to its large field-of-view, it only localizes them to $1^{\prime}$--$5^{\prime}$, which is not nearly adequate for finding optical/IR counterparts. Short {\em Chandra} exposures of IGR sources allow for a major advance in understanding the nature of these sources by providing sub-arcsecond positions, leading to unique optical/IR counterparts, as well as 0.3--10\,keV spectra that can be used to measure column densities and continuum shapes. After observations between 2003 and 2008, {\em INTEGRAL} had detected more than 700 hard X-ray sources \citep{bird10}. While that catalog includes nearly 400 IGR sources, the most up-to-date list\footnote{See http://irfu.cea.fr/Sap/IGR-Sources/.} includes more than 500 IGR entries. The largest group of identified sources are the Active Galactic Nuclei with $>$250 confirmed AGN detected \citep{bird10}. The largest Galactic source types include High-Mass X-ray Binaries (HMXBs), Low-Mass X-ray Binaries (LMXBs), and Cataclysmic Variables (CVs). A subset of these binaries are symbiotics with a white dwarf or neutron star accreting from a giant (i.e., luminosity class III) star. It has been somewhat surprising that the symbiotics with white dwarf accretors can produce hard X-ray emission \citep{kennea09,ratti10}. In addition, {\em INTEGRAL} has detected new as well as known supernova remnants (SNRs) and pulsar wind nebulae (PWNe). Unidentified IGR sources is also still a large group, and the most critical piece of information for determining the nature of these sources is to obtain a sub-arcsecond X-ray position. \begin{table*} \begin{center} \caption{{\em Chandra} Observations\label{tab:obs}} \begin{tabular}{cccclc} \hline \hline IGR Name & ObsID & $l$\footnote{Galactic longitude in degrees measured by {\em INTEGRAL}.} & $b$\footnote{Galactic latitude in degrees measured by {\em INTEGRAL}.} & Start Time (UT) & Exposure (s)\\ \hline \hline J01545+6437 & 12425 & 129.62 & +2.57 & 2011 Feb 7, 19.47 h & 4987\\ J03103+5706 & 12419 & 140.95 & --0.83 & 2010 Nov 23, 10.39 h & 4716\\ J04069+5042 & 12421 & 151.43 & --1.03 & 2010 Nov 25, 5.47 h & 4987\\ J06552--1146 & 12433 & 223.85 & --4.52 & 2010 Dec 12, 22.35 h & 4986\\ J09189--4418 & 12426 & 267.93 & +3.64 & 2011 Jul 22, 0.37 h & 5007\\ J12489--6243 & 12416 & 302.64 & +0.15 & 2010 Nov 25, 22.19 h & 4986\\ J13402--6428 & 12424 & 308.15 & --2.10 & 2011 Oct 6, 9.38 h & 4995\\ J15293--5609 & 12417 & 323.66 & +0.22 & 2011 Jun 22, 22.08 h & 5096\\ J15368--5102 & 12428 & 327.52 & +3.76 & 2011 Oct 9, 2.91 h & 4695\\ J15391--5307 & 12422 & 326.58 & +1.88 & 2011 Apr 14, 21.18 h & 4992\\ J15415--5029 & 12429 & 328.45 & +3.76 & 2011 Oct 9, 1.13 h & 4994\\ J16173--5023 & 12415 & 332.84 & +0.14 & 2011 Oct 11, 5.28 h & 4989\\ J16206--5253 & 12423 & 331.47 & --2.01 & 2011 Jul 5, 18.10 h & 4992\\ J16413--4046 & 12427 & 342.71 & +3.70 & 2011 May 9, 0.50 h & 4991\\ J16560--4958 & 12431 & 337.32 & --4.17 & 2011 May 23, 5.91 h & 4992\\ J21188+4901 & 12418 & 91.27 & --0.33 & 2011 Feb 28, 2.77 h & 4989\\ J21565+5948 & 12430 & 102.54 & +4.06 & 2011 May 30, 20.79 h & 4992\\ J22014+6034 & 12432 & 103.49 & +4.28 & 2011 Apr 16, 11.82 h & 4994\\ \hline \end{tabular} \end{center} \end{table*} Even the faintest {\em INTEGRAL} detections, which are near $\sim$$2\times 10^{-12}$\,erg\,cm$^{-2}$\,s$^{-1}$ (20--40\,keV), are usually easily detected by {\em Chandra} in a short exposure time. With our previous {\em Chandra} programs, we obtained 5\,ks snapshots of 46 IGR sources, and obtained 35 {\em Chandra} counterparts \citep{tomsick06,tomsick08a,tomsick09a}. Based on {\em Chandra} spectral information and optical/IR observations \citep[e.g.,][]{masetti7,chaty08,butler09}, definite or likely identifications have been obtained for 32 of our {\em Chandra} targets. These programs have all concentrated on low Galactic latitude targets as the {\em Chandra} positions are most critical for obtaining optical and IR counterparts in the crowded regions of the Galactic plane. Although our scientific goals have leaned toward finding Galactic sources such as obscured HMXBs \citep{mg03,fc04,walter06,cr12} or supergiant fast X-ray transients \citep[SFXTs,][]{negueruela06,smith06,pellizza06}, it has still been relatively common to find that the {\em INTEGRAL} sources are background AGN \citep[e.g.,][]{zct09}. In the following, we describe the results of {\em Chandra} observations of 18 IGR sources that were made between late-2010 and late-2011. The observations are described in \S\,2, and the search for X-ray counterparts, fits to the {\em Chandra} energy spectra, and multi-wavelength identifications are described in \S\,3. A discussion of the results for each source is provided in \S\,4, and \S\,5 includes a summary. | While we expect follow-up observations, especially optical and near-IR spectroscopy, to provide definite identifications of the nature of many of these sources, the {\em Chandra} observations have allowed us to determine the nature of some of the IGR sources as well as providing {\em Chandra} counterpart identifications for 14 of the 18 sources. A summary is provided in Table~\ref{tab:summary}, indicating that four of the IGR sources show definite or very strong evidence for being AGN. IGR~J01545+6437 and IGR~J21565+5948 both have redshift measurements, IGR J15391--5307 has a strong IR excess and radio emission, and IGR~J15415--5029 is identified with the Galaxy WKK98~5204 (2MASX~J15412638--5028233). IGR~J15415--5029 is one of the IGR sources for which we considered two {\em Chandra} candidate counterparts (9a and 9b), but we strongly suspect that 9a is the correct counterpart. Four other sources (IGR~J03103+5706, IGR~J09189--4418, IGR~J16413--4046, and IGR~J16560--4958) are likely AGN, and we label them as ``AGN?'' in Table~\ref{tab:summary}. For these sources, the evidence is extended optical or IR emission or an IR excess (or both features). The remaining six IGR sources for which we obtained one or two possible {\em Chandra} identifications may be Galactic sources. If the identification of IGR~J15293--5609 with CXOU~J152929.3--561213 is correct, then this is a definite Galactic source based on a parallax measurement that puts it distance at $1.56\pm 0.12$~kpc. Furthermore, based on the temperature and radius inferred from a blackbody fit to its optical/IR SED, it is likely a symbiotic binary with an early K-type giant donor, and its X-ray luminosity suggests a white dwarf accretor. IGR~J16206--5253 has {\em Chandra} and near-IR counterparts, and we argue from its SED that it is likely Galactic. IGR~J12489--6243 and IGR~J16173--5023 may be Galactic; however, there are still two possible {\em Chandra} counterparts for each source, and, in each case, there is some uncertainty about the nature of at least one of the counterparts. Once the nature of each of the {\em Chandra} counterparts is determined via follow-up measurements, it should become clear which counterpart is the most likely to be able to produce the hard X-ray emission detected with {\em INTEGRAL}. IGR~J15368--5102 has a {\em Chandra} counterpart, and IGR~J13402--6428 has two possible {\em Chandra} counterparts. However, the natures of these sources are still unclear. Figure~\ref{fig:summary} shows a plot of the $K_{s}$-magnitudes and the absorbed 2--10\,keV X-ray fluxes for the {\em Chandra} counterparts. The eight AGN or likely AGN fall within a range of $K_{s}$-magnitudes between 14.7 and 12.0 and within a range of X-ray fluxes between $2.5\times 10^{-13}$ and $5.7\times 10^{-12}$\,erg\,cm$^{-2}$\,s$^{-1}$. We compare these ranges to a previous study of AGN selected by 2--10\,keV fluxes measured by {\em ASCA}. \cite{watanabe04} show a similar $K_{s}$ vs. 2--10\,keV X-ray flux for $\sim$100 {\em ASCA} detected AGN, and plot lines of constant $\log{(f_{X}/f_{K_{s}})}$. They find that the {\em ASCA} AGN mostly fall between values of $\log{(f_{X}/f_{K_{s}})}$ 0 and 1, and Figure~\ref{fig:summary} shows that five of our AGN or likely AGN (2, 3, 12, 13, and 14) fall in that range. \cite{watanabe04} also included a sample of {\em Chandra} AGN in their study, and a significant fraction of those AGN have $\log{(f_{X}/f_{K_{s}})}$ values ranging from --1 to 0. Three of our AGN (1, 8, 9a) fall in this range. Thus, we conclude that the near-IR and X-ray fluxes that we find for our AGN are fairly typical. Galactic sources are not expected to be well-localized in a near-IR/X-ray flux plot, and Figure~\ref{fig:summary} shows that our Galactic or potentially Galactic sources are spread throughout the plot. The two sources with well-measured blackbody spectra are in the near-IR-bright/X-ray-faint corner of the diagram. The rest of these are relatively faint in the near-IR but have a wide range of 2--10\,keV X-ray fluxes. We have identified IGR~J15293--5609 (6) as a likely symbiotic binary with an early K-type giant donor and a white dwarf accretor. For 5a, 7, 10a, and 10b, higher angular resolution or deeper images are necessary to obtain optical/IR counterparts. For 4a, 4b, 5b, and 11, optical or near-IR spectra of the counterparts identified in this work are likely to lead to a determination of the source type. | 12 | 6 | 1206.1071 |
1206 | 1206.1592_arXiv.txt | We report the discovery of KELT-2Ab, a hot Jupiter transiting the bright (V=8.77) primary star of the HD 42176 binary system. The host is a slightly evolved late F-star likely in the very short-lived ``blue-hook'' stage of evolution, with $\teff=6148\pm48{\rm K}$, $\log{g}=4.030_{-0.026}^{+0.015}$ and $\feh=0.034\pm0.78$. The inferred stellar mass is $M_*=1.314_{-0.060}^{+0.063}$\msun\ and the star has a relatively large radius of $R_*=1.836_{-0.046}^{+0.066}$\rsun. The planet is a typical hot Jupiter with period $4.11379\pm0.00001$ days and a mass of $M_P=1.524\pm0.088$\mj\ and radius of $R_P=1.290_{-0.050}^{+0.064}$\rj. This is mildly inflated as compared to models of irradiated giant planets at the $\sim$4 Gyr age of the system. KELT-2A is the third brightest star with a transiting planet identified by ground-based transit surveys, and the ninth brightest star overall with a transiting planet. KELT-2Ab's mass and radius are unique among the subset of planets with $V<9$ host stars, and therefore increases the diversity of bright benchmark systems. We also measure the relative motion of KELT-2A and -2B over a baseline of 38 years, robustly demonstrating for the first time that the stars are bound. This allows us to infer that KELT-2B is an early K-dwarf. We hypothesize that through the eccentric Kozai mechanism KELT-2B may have emplaced KELT-2Ab in its current orbit. This scenario is potentially testable with Rossiter-McLaughlin measurements, which should have an amplitude of $\sim$44 m s$^{-1}$. | Individual giant planets transiting bright main sequence stars remain of prime scientific interest. While the multitude of hot Jupiters orbiting fainter (i.e., $V>9$) stars provides an opportunity to learn about the statistical properties of giant planets, it is the hot Jupiters around the bright stars which provide us with specific information about planetary interiors and atmospheres \citep[see, e.g.,][]{winn2010a}. Indeed, since their discovery, all of the transiting hot Jupiters orbiting bright ($V<9$) stars have been observed repeatedly from space and the ground for precisely this reason \citep{seager2010}. Since there are currently only five transiting giant planets in this magnitude range, discovering even one more substantially increases the opportunities for these important, detailed, follow-up observations. The Kilodegree Extremely Little Telescope (KELT) North transit survey is designed to find precisely these planets. KELT-North uses a small aperture telescope with a very wide field of view to observe a strip in declination that covers approximately 40\% of the Northern sky. The combination of a small light-collecting area and a wide field of view for KELT-North was chosen to efficiently survey all of the dwarfs stars between $8<V<10$ in our footprint. We specifically chose this magnitude range to cover the gap that exists between radial velocity surveys ($V\lesssim8.5$) and other transit surveys ($V\gtrsim10$). The KELT-North survey has been in operation since 2006, and we have been actively generating and vetting planet candidates since 2011 April. In this letter we describe the discovery and characterization of a hot Jupiter transiting the bright primary component of the HD 42176 binary system, which we hereafter refer to as the KELT-2 system. | The final planetary parameters for KELT-2Ab place it in a region of mass-radius parameter space that is already well populated by other hot Jupiters. What is noteworthy about this system is the brightness of the primary star, the primary's evolutionary state and the presence of a K2V common proper motion companion. At an apparent magnitude of V=8.77, KELT-2A is the ninth brightest star with a known transiting planet, and the third brightest discovered by a ground-based transit survey\footnote{According to the Extrasolar Planets Encyclopedia at the date of writing.}. This makes KELT-2Ab an excellent candidate for both space- and ground-based follow-up work. In terms of the bright ($V<9$) transiting planets, KELT-2Ab, at 1.52\mj, allows access to a region of the mass-radius diagram otherwise unprobed by the known bright systems. HD 209458b, HD 149026b and HD 189733b are all less than 1.15\mj, while HD 17156b and HAT-P-2b are both over 3\mj. WASP-33b does not have a well-constrained mass, and 55 Cnc e and Kepler-21b are both super-Earths. The star KELT-2A itself is in an interesting region of parameter space. Comparison to the Yonsei-Yale isochrones \citep{demarque2004} with our determined temperature, surface gravity, and metallicity suggests that KELT-2A has just left the main sequence and its convective core is in the process of halting hydrogen fusion and the star is transitioning to shell burning. This so-called `blue-hook' transition \citep[see, e.g.,][]{exter2010}, which occurs immediately prior to the star's rapid transition across the Hertzsprung gap to the base of the red giant branch, only lasts a few tens of millions of years. Assuming that the isochrones and our inferred stellar properties are correct in placing KELT-2A on the `blue-hook,' we thus find a remarkably precise system age of $3.968\pm0.010$ Gyr. The existence of the stellar companion KELT-2B raises the intriguing possibility that KELT-2Ab migrated inward to its present location through the eccentric Kozai mechanism \citep{lithwick2011}. If this were true, then the orbit of the planet is likely misaligned with the spin axis of KELT-2A. Interestingly, the effective temperature of KELT-2A (6151K) places this system near the proposed dividing line between cool aligned and hot misaligned planetary systems noted by \cite{winn2010}. Future Rossiter-McLaughlin measurements of the system's spin-orbit alignment should provide insight into the efficiency of any mechanisms that might align planets around the cooler stars. We would expect, from Equation (6) of \cite{gaudi2007}, the amplitude of the Rossiter-McLaughlin anomaly to be $\sim$44 m s$^{-1}$. Given the brightness and spectral type of KELT-2A, it is interesting to ask why this system was not observed by any of the RV surveys for exoplanets. In addition to KELT-2A being fainter than most of the targets for RV surveys, it may be that the RV surveys did not examine the KELT-2 system because it is listed as a binary in many of the available catalogs. | 12 | 6 | 1206.1592 |
1206 | 1206.6235.txt | The paper presents a sample of newly detected eclipsing binaries from the public $Kepler$ data. Orbits and fundamental parameters of 20 unknown eclipsing binaries were determined by modeling of their photometric data. Most of them are well-detached, high-eccentric binaries. We established that the target KID8552719 satisfied all widespread criteria for a planetary candidate. Fitting its light curve we obtained radius R$_{p}$=0.9 R$_{Nept}$, distance to the host star $a$=42.58 R$\sun$=0.198 AU and equilibrium temperatute T$_{p}$=489 K. These values imply a Neptune-size object out of the habitable zone of the host star. | Our searching of eclipsing binaries began in 2011 on the basis of the public data of quarters Q0-Q2 of the {\emph{Kepler} mission. Initially we used the data from the web-based NASA Exoplanet Archive http://exoplanetarchive.ipac.caltech.edu/\\ applications/ETSS/Kepler/ index.html. Further we used also the access to these data via ftp://130.167.252.39/pub/kepler/lightcurves/. It should be noted that we considered only the long cadence data because the short cadence ones were with bigger scatter The photometric data for each quarter and KID are calibrated for bias (dark level) and converted to fluxes. They are available in two forms: "RAW" and "CORR". Light curves for all observed stars (processed on a quarterly basis at NASA Ames Research Center) are placed in the archive. Users of {\emph{Kepler} data might work with both versions of the light curves. Particularly, everyone may use de-trended data or make de-trending himself. During 2011 the web-based NASA Exoplanet Archive contained not only stellar parameters (magnitudes, temperatures, etc.) but also the statistical parameters of their light curvesl "Light curve chi-squared", "Points in light curve beyond 5 $\sigma$" and "Fraction of points in light curve beyond 5 $\sigma$". We used the values of these statistical parameters to search for eclipsing binaries and managed to pick around thousand EB candidates. Unfortunately, the web-based NASA Exoplanet Archive do not contain yet the foregoing statistical parameters and the considered method for selection of EB candidates is not applicable. The next step was visual examination of our EB candidates. It revealed that most of them are false positives and only around hundred are true EB. Besides the foregoing approach we made visual inspection of all light curves of one third of the Kepler data (until folder 0040 from the downloaded ftp data). As a result of these two ways of searching we gathered candidates that were out of the published catalogs of $Kepler$ eclipsing binaries (Prsa' catalog version 2.0). Further we checked if some of these candidates were known variables (to the middle of May 2012). For this aim we used: (i) Lists of Planetary Candidates, False Positives and Eclipsing Binaries from the site http://archive.stsci.edu/kepler/results.html; (ii) The AAVSO database of variable stars (International Variable Star Index) (iii) Databases SIMBAD and ADS. Several candidates turned out known variables and were removed from our sample. Thus we made the first version of our list of EB candidates. In this paper we present part of the newly detected eclipsing stars satisfying the additional condition to have at least 3 eclipses from different type. During the visual inspection of the light curves of the \emph{Kepler} database we established that they need additional preliminary "manual" treatment before detailed analysis because the maximum fluxes as well as the light amplitudes for many stars differ from one quarter (even subset) to another. Moreover, the "CORR" data are not always properly de-trended. That is why we decided to use the RAW data of our targets and to made own normalization and de-trending of the data. For this aim the data from each subset were reproduced by spline function and normalized separately. At the beginning of 2012 when the quarters Q3-Q6 became free-accessed we used the new data for more precise determination of the ephemeres of our targets. The last examination of our targets consisted in checking for possible mis-classified of variability of some target as a result of contamination of nearby known binary with the same period. This effect is due to the high star density near the galactic plane and to the fact that the \emph{Kepler} pixels projected onto 4 arcsec on the sky. Thus, for instance, we found that the unknown variable star KID3098197 (detected by us) and the known EB star KID3098194 (Paper I) at distance 4 arcsec have the same period and type of variability (Fig. 1). It is difficulty to say which of these stars is the true variable because the contamination depends on the relative magnitude of the stars in question and their separation on the sky. As a result of the examination for contamination effect (into area with radius 60 arcsec around our candidates) several candidates as KID3098197 were excluded from our list. \begin{figure}[!htb] \begin{center} \centering{\epsfig{file=KICneihbour3.eps, width=0.95\textwidth}} \caption[]{Illustration of contamination effect by the stars KID3098194 and KID3098197} \label{1} \end{center} \end{figure} | Applying own searching method for eclipsing binaries to the public $Kepler$ database Q0-Q2 we detected candidates that were out of the published catalogs of the $Kepler$ eclipsing stars. This paper presents the results of our light curve solutions of 20 newly detected eclipsing binaries and the values of their orbital and fundamental parameters. Most of them turned out binaries with high-eccentric orbits and consequently present appropriate targets for study of apsidal motion and the internal structure of the component stars. Some of our discoveries have interesting peculiarities. We confirmed the conclusion of Huang, Bakos $\&$ Hartman (2012) that KID8552719 satisfied the criteria for a planetary candidate. The estimations for its physical parameters revealed the Neptune-size object out of the habitable zone of the host star. Hence, the softwares do are indispensable for an automatic processing and analysis of a huge data-sets of the space missions as \emph{Kepler} but the obtained results in this paper revealed that it is deserved to examine such data-sets "manually/visually" in order to compensate the omission of some interesting objects and events, available in the database but not registered by the software. At first glance this time-consuming work seems boring and requires patience and attention. But it is a true pleasure to see the wonderful light curves of the \emph{Kepler} mission! | 12 | 6 | 1206.6235 |
1206 | 1206.5558_arXiv.txt | {It is clear that optical selection effects have distorted the ``true" GRB redshift distribution to its presently observed biased distribution. We constrain a statistically optimal model that implies GRB host galaxy dust extinction could account for up to 40\% of missing optical afterglows and redshifts in $z = 0-3$, but the bias is negligible at very high-$z$. The limiting sensitivity of the telescopes, and the time to acquire spectroscopic/photometric redshifts, are significant sources of bias for the very high-$z$ sample. We caution on constraining star formation rate and luminosity evolution using the GRB redshift distribution without accounting for these selection effects. } \FullConference{Gamma-Ray Bursts 2012 Conference -GRB2012,\\ May 07-11, 2012\\ Munich, Germany} \begin{document} | Figure \ref{fig_pdf_results} plots the expected redshift distribution including these completeness functions. The optimal model based on a Kolmogorov Smirnov (KS) probability $>95\%$ is one that includes selection effects. {\bf Dust extinction} produces a $(30-40)$\% reduction in redshifts in $z=0-3$. The fraction of missed redshifts from the {\bf reshift desert} is $<25$\% in $z=1.5-2.5$. At high-$z$ the {\bf OA Malmquist bias} results in 20\% of redshifts missed in $z=0-5$ and increases to 50\% missed out to $z=10$. | 12 | 6 | 1206.5558 |
|
1206 | 1206.0182_arXiv.txt | We analysed all archival RXTE observations of the neutron-star low-mass X-ray binary 4U 1636--53 up to May 2010. In 528 out of 1280 observations we detected kilohertz quasi-periodic oscillations (kHz QPOs), with $\sim 65$\% of these detections corresponding to the so-called lower kHz QPO. Using this QPO we measured, for the first time, the rate at which the QPO frequency changes as a function of QPO frequency. For this we used the spread of the QPO frequency over groups of 10 consecutive measurements, sampling timescales between 320 and 1600 s, and the time derivative of the QPO frequency, $\dot \nu_{\rm QPO}$, over timescales of 32 to 160 s. We found that: (i) Both the QPO-frequency spread and $\dot \nu_{\rm QPO}$ decrease by a factor $\sim 3$ as the QPO frequency increases. (ii) The average value of $\dot \nu_{\rm QPO}$ decreases by a factor of $\sim 2$ as the timescale over which the derivative is measured increases from less than 64 s to 160 s. (iii) The relation between the absolute value of $\dot \nu_{\rm QPO}$ and the QPO frequency is consistent with being the same both for the positive and negative QPO-frequency derivative. We show that, if either the lower or the upper kHz QPO reflects the Keplerian frequency at the inner edge of the accretion disc, these results support a scenario in which the inner part of the accretion disc is truncated at a radius that is set by the combined effect of viscosity and radiation drag. | Kilohertz quasi-periodic oscillations (kHz QPOs) are the fastest variability so far observed in neutron star (NS) low-mass X-ray binary (LMXB) systems. These oscillations were first detected (\citealt{v1996a}; \citealt{st1996}) shortly after the launch of the Rossi X-ray Timing Explorer (hereafter RXTE; see \citealt{b1993}; \citealt{ja2006}). Since then, kHz QPOs have been observed in more than 30 neutron star LMXBs (see, e.g., \citealt{v2005}). KHz QPOs are often observed in pairs, usually called lower and upper kHz QPO, with frequencies ranging from a few hundred Hz to more than 1 kilohertz (e.g., \citealt{v2005}, for an extensive review). The fact that the timescale of these oscillations corresponds to the dynamical timescale very close to the NS, makes kHz QPOs potentially one of the few tools nowadays available to directly measure strong-field gravitational effects. In the past 16 years, several models have been proposed to explain these oscillations, with emphasis in trying to explain the frequency of the QPOs. Almost all models associate the frequency of these oscillations with characteristic frequencies in a geometrically thin accretion disc (e.g., \citealt{mlp1998}; \citealt{sv1998}). Systematic studies of other properties of kHz QPOs, such as the coherence and the amplitude of the oscillation, are also available for several LMXBs (\citealt{j2000}; \citealt{van2000}; \citealt{ds2001}; \citealt{mvf2001}; \citealt{h2002}; \citealt{van2002}; \citealt{dmv2003}; \citealt{van2003}; \citealt{b2005a}; \citealt{b2005b}; \citealt{bom2006}; \citealt{a2008}; \citealt{sa2010}), but not many models explain these other properties (see \citealt*{mlp1998}; \citealt*{bom2006}). Some sources show a characteristic dependence of the coherence (defined as $Q = \nu/\Delta \nu$, where $\nu$ and $\Delta\nu$ are the centroid frequency and the full width at half-maximum of the kHz QPO, respectively) and the amplitude of those oscillations upon the centroid frequency of the oscillation: Both the coherence and the amplitude of the lower kHz QPO in these sources increase as the frequency increases, and then quickly decrease as the QPO frequency continues to increase. \citet*{bom2006} proposed that the drop of the coherence $Q$ at high frequencies is a direct consequence of the disc approaching the innermost circular stable orbit (ISCO), where the inner edge of the disc gets shattered by the gravitational force of the neutron star. If this interpretation is correct, the drop of the coherence and rms amplitude of the lower kHz QPO at high frequencies would confirm the existence of the ISCO, and it would provide a direct measurement of the mass of the neutron star \citep*{mlp1998}. However, other non space-time-related interpretations have been proposed to explain the properties of QPOs; e.g., \citet{m2006} suggested that the drop of $Q$ and rms amplitude in individual sources might be related (at least in part) to changes in the properties of the accretion flow in these systems. This could explain the significantly different properties of the kHz QPOs, respectively, in the high and low-luminosity phase of the outburst of the LMXB XTE J1701-462 \citep{sa2010}. Here we investigate the properties of the kHz QPOs for the NS system 4U 1636-53 by scanning the whole RXTE archive until May 2010. In Section 2 we describe the observations and the data analysis, and in Section 3 we present our results. In Section 4 we discuss our results in the context of current ideas concerning the origin of the kHz QPOs in LMXBs. \section[]{OBSERVATIONS AND DATA ANALYSIS} \label{data} We analysed all archival observations of the LMXB 4U 1636--53 obtained with the Proportional Counter Array (PCA) on board of RXTE up to May 2010. We used a total of 1280 RXTE observations, which correspond to an exposure time of $\sim$ 3.5 Ms. During these observations the source showed a large number of type-I X-ray bursts that we excluded from our analysis (see, \citealt{z2011}). Except for four observations, we used event-mode data with 125$\mu s$ time resolution and 64 channels covering the full PCA energy band to produce Fourier power density spectra (PDS), setting the time resolution to 4096 points per second, corresponding to a Nyquist frequency of 2048 Hz. In the four observations in which the event-mode data were not available, we used a combination of binned modes that allowed us to reach the same Nyquist frequency in the PDS. In all cases we selected photons with energies below $\sim$46 keV (absolute channel 101), with the exception of observation 10072-03-01-00 where we selected photons between channels 24 and 139, and observations 10072-03-02-00 and 10088-01-08-010 where we used the full energy band, because in these cases it was not possible to have simultaneously the same Nyquist frequency and energy selection as in the other observations. We created Leahy-normalised PDS for each 16-seconds data segment. We removed detector drop-outs, but no background subtraction or dead-time correction were applied before calculating the PDS. \\ We averaged all 16-s spectra within each observation, and we searched for kHz QPOs by fitting the averaged PDS in the frequency range 200-1500 Hz using a constant to model the Poisson noise and one or two Lorentzians to model the QPOs. From each fit, we estimated the significance of the Lorentzian by dividing its normalisation by the negative $1\sigma$ error. We considered only features with this ratio larger then 3, and coherence Q larger than 2.\\ We calculated the X-ray colours and intensity of the source using the Standard 2 data (16-s time resolution and 129 channels covering the entire PCA energy band). We defined soft and hard colour as the count rate ratio in the energy bands 3.5-6.0 keV / 2.0-3.5 keV and 9.7-16.0 keV / 6.0-9.7 keV, respectively. The intensity of the source is defined as the count rate in the energy band 2.0-16.0 keV. To obtain the exact count rate in each of these energy bands we linearly interpolated in channel space, since the energy boundaries of each channel change slightly with time. To correct for gain changes and differences in the effective area between the proportional counter units (PCUs) as well as differences due to changes in the channel to energy conversion of the PCUs as a function of time, we normalised the colours and intensity to those of the Crab Nebula obtained close in time to our observations (see \citealt{k2004} and \citealt{a2008} for details). Finally we averaged the normalised colours and intensities per PCU for the full observation using all available PCUs.\\ \subsection{KHz QPOs Identification} \label{QPOiden} We detected kHz QPOs in 528 out of 1280 observations analysed. Among the observed QPOs, $\sim$84\% where detected with a significance larger then 3.5$\sigma$, $\sim$72\% with a significance larger then 4$\sigma$, and $\sim$53\% with a significance larger then 5$\sigma$. Even though $\sim$16\% of the QPOs could be considered only marginally significant (between 3 and 3.5$\sigma$), all these QPOs follow the frequency-hardness correlation described below, which adds confidence to the detections. Only 26 observations with QPOs showed two simultaneous high-frequency oscillations. In those cases the QPO identification is trivial because lower and upper kHz QPOs were both present. For all the other observations where we only detected a single QPO peak we had to apply a different method to identify the peak as either the lower or the upper kHz QPO. In Figure \ref{hard_freq} we plot, for each observation, the centroid frequency of all the kHz QPOs as a function of the hard colour of the source. From the plot it is apparent that, for frequencies below 1000 Hz, the data follow two separate branches. The first branch extends between QPO frequencies of $\sim$ 550 Hz and $\sim$ 950 Hz, with more or less constant hard colour at around 0.65, except for a small deviation at low frequencies where the hard colour increases slightly. The second branch shows a clear anti-correlation between frequency and hard colour. This has been already reported for this source (see \citealt{b2007}) as well as for other neutron-star systems (e.g., \citealt{m1999}; \citealt{mv1999}) and can be used to identify the two different kHz QPOs: Lower kHz QPO at low hard colour, and upper kHz QPO at high hard colour. In Figure \ref{hard_freq} we also show the kHz QPOs detected simultaneously per observation: Red-empty triangles and blue-empty circles represent, respectively, the lower and the upper kHz QPOs. Those points confirm the correlation previously mentioned between frequency, hard colour and type of QPO. Interestingly, by checking the distribution of the blue circles in Figure \ref{hard_freq} we can resolve the ambiguity regarding the identification of the QPOs in the region above 1000 Hz, in which the two branches approach each other. To summarise, we detected a total of 357 lower kHz QPOs and 197 upper kHz QPOs. Compared to \citealt{b2007}, we find that the branch corresponding to the upper kHz QPO extends to higher frequencies at more or less constant hard colour, while the branch corresponding to the lower kHz QPO extends to lower frequencies at higher hard colour. \begin{figure} \begin{center} \includegraphics[width=80mm]{figure1.eps} \caption{Centroid frequency of the kHz QPOs detected in 4U 1636-53 as a function of hard colour. Grey bullets represent the QPOs in observations in which only one kHz QPO is observed. Red triangles and blue circles represent QPOs identified, respectively, as the lower and the upper kHz QPO in observations in which two simultaneous kHz QPOs are detected. For clarity, in the inset we plot only the QPO frequencies of the observations in which two simultaneous kHz QPOs were detected.} \label{hard_freq} \end{center} \end{figure} \subsection{Variations of the kHz QPO frequency} \label{sp} It is well known that the frequency of the kHz QPOs can change over tens of Hz in time intervals of a few hundred seconds \citep[e.g.,][]{b1996}, and this can artificially broaden the QPO in the averaged power spectrum of long observations. In order to study frequency variations of the kHz QPOs, we tried to detect kHz QPOs on timescales shorter than the full observation. We found that in observations with only the upper kHz QPO we cannot detect the QPO significantly on timescales shorter than the full observation, whereas we detected the lower kHz QPO on timescales between 32 s and 160 s in about 50\% of the observations that showed kHz QPOs, all at significance levels larger than 4$\sigma$. For these observations we created dynamical power density spectra \citep[see, e.g., Fig.2 in][]{b1996}, and we extracted the frequency evolution as a function of time (hereafter frequency profile). To reduce the influence of the frequency error on the analysis of the frequency variation, we applied a Savitzky--Golay filter (\citealt{sk1964}; \citealt{nr}) to smooth the frequency profile. This method locally fits a polynomial function on a series of values (in our case we used a 4th order polynomial on 6\footnote{We varied the number of overlapping points from 4 to 10, but within errors this does not affect the time derivative of the QPO frequency.} consecutive points), and it provides the first time derivative of the QPO frequency (hereafter frequency derivative) for each measurement. For each observation we then calculated how much the centroid frequency moved with time by measuring the sample standard deviation (hereafter the spread) using groups of 10 consecutive measurements of the frequency profile. We also used the frequency derivative to evaluate the rate at which the centroid frequency changes on certain timescales. \section[]{RESULTS} \label{res} In Figure \ref{spread} we show the spread of the lower kHz QPO frequency as a function of the QPO frequency measured over timescales of 320 and 1600 s. Each point represents the average spread within frequency bins of 5 to 40 Hz, depending on the number of measurements at different frequencies. Notice that the rebin has been applied under the assumption that the spread only depends on frequency. Figure \ref{spread} shows that the frequency spread decreases from $\sim$ 7 Hz to $\sim$ 2 Hz as the QPO frequency increases from $\sim$ 650 Hz up to $\sim 910$ Hz. Above $\sim 910$ Hz, the data show a marginal increase of the spread to $\sim$ 4 Hz as the frequency increases further up to $\sim$ 930 Hz. The apparent rise of the spread at high frequency is suggested by two points only, which are the average of 13 and 12 measurements, respectively. We fitted the data both with a line and a broken line function, and we carried out an F-test to compare both fits. The F-test probability is $2.2 \times 10^{-2}$, indicating that the fit with a broken line is marginally better than a fit with line. We cannot rule out the possibility that this increase is real, but as we discuss below, it could also be (partly) due to a larger contribution of the statistical errors to the QPO frequency in that frequency range. To take into account the contribution of the statistical errors from the measurements of the frequency on the spread, we subtracted in quadrature the 1-sigma errors from the observed spread (see \citealt{vau2003}, section 6.1, Eq. 8). We found that the errors of the QPO frequency contribute about 10\% of the observed spread at around 650 Hz, whereas this contribution decreases down to 5\% at 850 Hz, and then increases up to 40\% at $\sim$ 930 Hz. Although this correction slightly flattens the relation, the overall trend does not change significantly, except that the apparent increase of the spread at the high-frequency end becomes less significant. The accuracy with which one measures the centroid frequency of the QPO depends on the amplitude and the width of the oscillation (e.g., \citealt{v1997}), which in turn change with QPO frequency (e.g., \citealt*{bom2006}). In order to verify whether this affects the relation between spread and frequency, we made simulations where we added normally distributed noise to the frequency profile, with standard deviation ranging from 1 to 10 Hz. As expected, the extra noise component tends to flatten the relation, but errors of up to 10 Hz (rms) in the determination of the QPO frequency do not significantly change the trend shown in Figure \ref{spread}. \begin{figure} \begin{center} \includegraphics[width=80mm]{figure2.eps} \caption{Measurements of the frequency spread of the lower kHz QPO in 4U 1636--53 as a function of the QPO centroid frequency. Each point is the average of the QPO-frequency spread at that frequency, corresponding to timescales between 320 s and 1600 s. See section \ref{res} for more details.} \label{spread} \end{center} \end{figure} Figure \ref{speed64} shows the absolute value of the frequency derivative of the lower kHz QPO as a function of the QPO frequency. This plot is based on measurements of the QPO frequency on timescales of 64 s or less. Each point represents the average value of the frequency derivative within a frequency bin from 5 to 30 Hz. Similar to what happens with the spread, the frequency derivative decreases as the frequency of the QPO increases, but unlike Figure \ref{spread}, in this case there is no indication of an increase of the frequency derivative above $\sim$ 910 Hz. \begin{figure} \begin{center} \includegraphics[width=80mm]{figure3.eps} \caption{Absolute value of the frequency derivative of the lower kHz QPO frequency in 4U 1636--53 as a function of the QPO frequency. Each point is the average value (within a frequency bin) of the frequency derivative of the QPO frequency detected on timescales of 64 s or less. See section \ref{sp} for more details.} \label{speed64} \end{center} \end{figure} In Figure \ref{speed_allT} we compare the absolute value of the frequency derivative of the lower kHz QPOs detected over three different timescales. Black points represent the frequency derivative of the QPOs measured on timescales of 64 s or less (the same as Figure \ref{speed64}), red points are for timescales between 64 and 160 s, and green points are for timescales of 160 s. Interestingly, for timescales longer than 64 s the frequency derivative initially stays more or less constant as the QPO frequency increases, and at around 800 Hz it decreases following more or less the same trend seen on the shortest timescale. We investigated whether the difference between trends in Figure \ref{speed_allT} is only caused by the different timescales. To do that we took the observations where we detected kHz QPOs in segments of data shorter than 64 s (black points in Figure \ref{speed_allT}), we created PDS increasing the length of the data segments up to 160 s and we searched again for kHz QPOs. We then recalculated the frequency derivative and we compared it with the green points in Figure \ref{speed_allT}. We found that the two distributions of points are consistent with each other, which shows that the difference between the three groups of points in Figure \ref{speed_allT} is likely due to the different timescales over which we measured the QPO frequency. Finally, in Figure \ref{speed_pos_neg} we show, respectively, the absolute value of the negative (frequency decrease) and positive (frequency increase) frequency derivative for kHz QPOs detected on timescales of up to 64 s. We fitted both relations with a linear function. We found that, within the errors, there is no significant difference between the trend of the positive and negative QPO-frequency derivative with QPO frequency. \begin{figure} \begin{center} \includegraphics[width=80mm]{figure4.eps} \caption{Absolute frequency derivative values of the lower kHz QPO in 4U 1636--53 for 3 different timescales. Black points are measurements for QPOs detected on intervals shorter than 64 s. Red points represent detections between 64 s and 160 s, while green points represent measurements of the QPO on intervals of 160 seconds.} \label{speed_allT} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=80mm]{figure5.eps} \caption{Comparison between absolute values of the positive (red points) and negative frequency derivative (black points) for the lower kHz QPOs in 4U 1636--53 detected on timescales shorter than 64 s.} \label{speed_pos_neg} \end{center} \end{figure} | We detected kHz QPOs in 528 out of 1280 RXTE observations of the LMXB 4U 1636--53 (all RXTE data available of this source up to May 2010). 26 of these 528 observations showed simultaneously the lower and the upper kHz QPO, whereas in the remaining 502 observations we detected a single kHz QPO that we identified using the frequency vs. hard colour diagram (see section \ref{QPOiden}). Using measurements of the lower kHz QPO on timescales between 32 and 160 s we measured, for the first time, the rate at which the QPO frequency changes as a function of the QPO frequency itself. For this we used both the spread of the QPO frequency over groups of 10 consecutive measurements, sampling timescales between 320 and 1600 s, and the time derivative of the QPO frequency over timescales of 32 to 160 s. Both the QPO-frequency spread and QPO-frequency derivative decrease by a factor of $\sim$ 3 as the QPO centroid frequency increases from $\sim$ 600 to $\sim$ 900 Hz. Above $\sim$ 910 Hz the QPO-frequency spread appears to increase again, albeit the evidence for this is marginal. We found that the relation between QPO frequency, $\nu_{\rm QPO}$, and QPO-frequency derivative, $\dot{\nu}_{\rm QPO}$, depends upon the timescale over which the QPO frequency is measured: When $\nu_{\rm QPO}$ is between 600 and $\sim$ 770 Hz, $\dot{\nu}_{\rm QPO}$ decreases by a factor of $\sim$ 2 as the timescale over which we measure the frequency derivative increases from 64 s or less up to 160 s. Finally, we found that the positive and negative values of $\dot \nu_{\rm QPO}$ are consistently distributed in the whole frequency range (see Figure \ref{speed_pos_neg}), even at high frequencies where the accretion disc should approach the ISCO. \citet*{bom2006} found that in 4U 1636--53, the coherence of the lower kHz QPO increases from $\sim$10 to $\sim$220 as the frequency of the QPO increases from $\sim$600 to $\sim$850 Hz, and then it decreases rather abruptly as the frequency continues increasing up to 950 Hz\footnote{Data in Figure \ref{spread} to \ref{speed_pos_neg} show no unusual change in the trend of the QPO frequency spread or the time derivative at around 850 Hz, where \citet*{bom2006} see the sudden drop of the coherence of the lower kHz QPO in 4U 1636--53. Similarly, there are no unusual features at frequencies corresponding to small integer ratios of the kHz QPO frequencies in the spread and $\dot \nu_{\rm QPO}$ relations (\citealt{a2003}; \citealt{bmh2005}).}. \citet*{bom2006} were able to model the behaviour of the coherence of the kHz QPO in terms of the lifetime of blobs of matter that move in quasi-Keplerian orbits at the inner edge of a geometrically-thin disc, the width of the disc annulus where these blobs produce the kHz QPOs, and the advection speed of the gas in the disc. \citet*{bom2006} conclude that the abrupt drop of the QPO coherence at about 850 Hz is due to the presence of the ISCO around the neutron star in this and other systems. Comparing our results with Figure 8 in \citet{bom2006}, we find that the variation of the QPO frequency on timescales between 320 and 1600 seconds contributes significantly to the QPO width in the frequency range between $\sim$750 and $\sim$870 Hz, whereas it does not contribute to the increase of the QPO width at higher QPO frequencies. Previous results show that variations on shorter timescales can have a large impact on the width of QPO. E.g., \citet{ykj2001}, showed that in the neutron star system Sco X-1 the frequency of the upper kHz QPO changes by ~20 Hz on timescales of ~1/6 Hz, which they relate to flux variation on these timescales. The trend of the QPO spread in Figure \ref{spread} shows a marginal increase of the spread at around 915 Hz, which is close to the frequency at which \citet*{b2005c} identify the signature of the ISCO on the lower kHz QPO in 4U 1636--53. If this apparent rise was real, it could be related to the idea that the high-frequency drop of the QPO coherence is a signature of the disc reaching the ISCO. Most of the models so far proposed assume that the kHz QPOs are produced in an optically-thick geometrically-thin accretion disc, and that the frequency of either the lower or the upper kHz QPO represents the Keplerian frequency at a given radius in the disc (e.g., \citealt*{mlp1998}; \citealt{sv1998};\citealt{ot1999}). For instance, in the sonic-point model \citep*{mlp1998} the frequency of the upper kHz QPO corresponds to the radius at which the gas starts spiralling inward at a supersonic radial velocity, and changes of the QPO frequency reflect changes of mass accretion rate (ram pressure and radiation drag). In other models (e.g., \citealt{sv1998}; \citealt{ot1999}) the frequency of the QPOs also depends upon the radius of the disc where the QPOs are produced, although the mechanism by which this radius changes is not specified. In those models mass accretion rate is also likely responsible for changes of the radius in the disc where the QPO is produced. \subsection{Comparison with standard disc theory} We can compare our measurements of $\dot{\nu}_{\rm QPO}$ of the lower kHz QPO with what is expected in the case of a thin disc \citep{ss1973} if $\nu_{\rm QPO}$ is the Keplerian frequency very close to the inner edge of the disc. We will further assume that the disc is truncated by an unspecified mechanism. According to the standard disc model \citep{ss1973}, matter in the disc moves radially with a speed given by: \begin{equation} v_{r}=0.98\, \alpha_{s}^{4/5}\, \dot{M}^{3/10}\, M_{\star}^{-1/4}\,R^{-1/4}\,f^{-14/5}\,\text{km/s},\\ \end{equation} with $f=[1-(R_{\star}/R)^{1/2}]^{1/4}$, where $\alpha=0.1\alpha_{s}$ is the viscosity parameter, $\dot{M}$ is the mass accretion rate in units of $2.6\times10^{17}$ g/s, $M_{\star}=M/1.8M_{\odot}$, $R$ is the radius of the inner edge of the disc in km, and $R_{\star}$ is the neutron star radius in units of 14 km, respectively\footnote{We assumed an average luminosity value for 4U 1636--53 of about 0.1 $L_{\rm Edd}$ \citep{b3} for a distance of 5.5 Kpc \citep{vpw1995}, where $L_{\rm Edd}$ is the Eddington luminosity for a 1.8-solar masses neutron star.}. This is the speed at which matter very close to the inner edge of the accretion disc will move inwards one the disc has been truncated. Combining this expression with the Keplerian formula, we find that the derivative of the Keplerian frequency can be expressed as:\\ \begin{equation} \label{nu_dot} |\dot{\nu}| \approx 300\,\alpha_{s}^{4/5}\, \dot{M}^{3/10}\, M_{\star}^{-2/3}\,\nu_{kHz}^{11/6}\,F(\nu_{kHz})\,\text{Hz/s}, \end{equation} where $\nu_{kHz}$ is the Keplerian frequency in units of 1000 Hz, and $F(\nu_{kHz})$ is the factor $f^{-14/5}$ expressed as a function of $\nu_{kHz}$, and normalised by the same factor calculated for a frequency of 1000 Hz. Assuming a 1.8 $M_{\sun}$ neutron star, the function $F(\nu_{kHz})$ increases from $\sim0.6$ to $\sim0.9$ as the frequency increases from 650 to 930 Hz. From this equation it is apparent that the time derivative of the Keplerian frequency very close to the inner edge of a thin disc increases as the Keplerian frequency increases, contrary to what we observe in the case of the lower kHz QPO in 4U 1636--53. Not only that, but for nominal values of the parameters, the time derivative of the Keplerian frequency in the frequency range between 650 and 950 Hz, is about 4 orders of magnitude larger than what we observe for the lower kHz QPO in 4U 1636--53. As it is apparent from eq. \eqref{nu_dot}, this difference cannot be accounted for by changing the mass and the radius of the NS. One way to reconcile the observations with the expectations from a thin disc is if the viscosity parameter $\alpha$ is $\sim10^{-5}$ (cf. \citealt{b1997}), although this would not revert the trend of $\dot \nu_{\rm QPO}$ with the QPO frequency. The $\dot \nu_{\rm QPO}$ vs. $\nu_{\rm QPO}$ trend from our analysis could be in principle reproduced by eq.\eqref{nu_dot} if the product $\alpha \times \dot{M}$ depended upon the accretion-disc radius. If we assume for simplicity that $\alpha$ is constant with radius, the trend in eq.\eqref{nu_dot} would be reversed if the mass accreted from the secondary flowed out of the disc at a rate that increases as the radius decreases. In order to fit the observations, in 4U 1636--53 the amount of accreted mass that crosses through the inner edge of the disc must be $\sim 5$\,\% of the total accreted mass from the secondary. In other words, about 95\,\% of the mass accreted from the secondary should not flow onto the neutron star via de accretion disc. Although this scenario cannot be completely ruled out, it is very unlikely because such a process should leave some trace in the spectral properties of the system, e.g., extra low-energy absorption, that has not been observed so far. The majority of the models of the kHz QPOs propose that it is the {\em upper} kHz QPO that is Keplerian (e.g., \citealt*{mlp1998}, \citealt{sv1998}). Since the upper kHz QPO in 4U 1636--53 is much broader and weaker than the lower one (e.g.; \citealt*{b2005b}), the upper kHz QPO is in general less significant than the lower one when both are measured over the same timescale. We therefore were unable to recover the frequency profile of the upper kHz QPO on the same timescales as for the lower kHz QPO to compare the time derivative of the upper kHz QPO with that of the Keplerian frequency in a thin disc. However, \citet{bmh2005} showed that in 4U 1636--53 the frequency of the lower and upper kHz QPOs, $\nu_l$ and $\nu_u$, respectively, follow a linear relation $\nu_u = 0.673\times\nu_l + 539$ Hz. If this relation still holds on short timescales, we can combine it with our measurements of the lower kHz QPO and estimate the time derivative of the upper kHz QPO frequency. We find that, similar to the case of the lower kHz QPO described above, the time derivative of the frequency of the upper kHz QPO decreases as the frequency of the upper kHz QPO increases, which is again at variance with what is expected from the thin disc model. From the above we conclude that changes of the QPO frequency, at least for the case of 4U 1636--53, are not compatible with the dynamics of the inner edge of a standard $\alpha$-disc. \subsection{Standard disc theory plus radiation drag} \label{miller} The positive and negative derivatives of the QPO frequency appear to follow the same trend as a function of the QPO frequency (see section \ref{res} and Figure \ref{speed_pos_neg}). This suggests that there is a coupling between the mechanisms that drive the inner edge of the disc inwards and outwards. In this subsection we explore the case in which the accretion disc is truncated at the so-called sonic-point radius, the position in the disc where the radial velocity of the inflowing gas changes from subsonic to supersonic \citep*{mlp1998}. In this scenario, mass accretion rate through the accretion disc pushes the inner edge of the disc inwards in a similar way to the one we described in the previous subsection; but in this case, the radiation produced by the mass that eventually accretes onto the neutron-star surface removes angular momentum from the disc. The position of the inner radius of the accretion disc is therefore set by the interplay between radiation drag and viscosity. Using general relativistic calculations of the gas dynamics and radiation transport in the inner edge of the accretion disc, \citet*{mlp1998} found that the sonic radius can be written approximately as: \begin{equation} \label{radius_mill} R_{\textsl{sr}} \approx R+5\left(\frac{\dot{M_i}}{0.01\dot{M_E}}\right)^{-1} \left(\frac{R}{10\text{km}}\right)\left(\frac{h/r}{0.1}\right)\left(\frac{v^{r}}{10^{-5}\textsl{c}}\right), \end{equation} where $R$ is the NS radius in kilometres, $\dot M_i$ is the mass accretion rate through the disc, $\dot M_{E}$ is the Eddington mass accretion rate, $h$ is the thickness of the disc at a radial distance $r$ in the disc, and $v^r$ is the radial velocity of the gas in the accretion disc. We can therefore write: \begin{equation} |\dot{\nu}_{\textsl{sr}}| = \frac{3}{4\pi} \left(\textsl{GM}\right)^{1/2}\,\textsl{R}^{-5/2}_{\textsl{sr}}\,|\dot{\textsl{R}}_{\textsl{sr}}|, \end{equation} with \begin{equation} \label{nu_sr} |\dot{\textsl{R}}_{\textsl{sr}}|\approx 0.05\,\dot{M}_{E}\left(\frac{\dot{M_i}}{0.01\dot{M}_{E}}\right)^{-2}\ddot{M_i} \left(\frac{R}{10\text{km}}\right)\left(\frac{h/r}{0.1}\right)\left(\frac{v^{r}}{10^{-5}\text{c}}\right), \end{equation} where $\ddot M_i$ is the time derivative of the mass accretion rate through the disc. We can now in principle use equations (3) to (5) to calculate $\dot \nu_{\rm sr}$ vs. $\nu_{\rm sr}$, given the NS mass and radius. For that we also need to specify $h/r$, $v^{r}$ and $\ddot M_i$. In systems in which accretion is the main source of radiation, the bolometric luminosity is proportional to the total mass accretion rate. To the extent that in LMXBs the X-ray intensity, $I_{\rm X}$, is a good measure of the bolometric luminosity (e.g., \citealt{b3}), $\ddot M_i \propto \dot I_{\rm X}$. Putting all this together, and assuming that $h/r$ and $v^r$ are constant, we find that $\dot \nu_{\rm sr} \propto \nu_{\rm sr}^{5/3}\, \dot I_X / I_X^2$. Using the observed variations of $I_{\rm X}$ with time we find that $\dot\nu_{\rm sr}$ decreases by a factor $\sim 2$ as $\nu_{\rm sr}$ increases from $\sim 650$ to $\sim 900$ Hz, which is comparable to what we found in 4U 1636--53 (see Figure \ref{speed64}). However, it is well known that, while on timescales of hours kHz QPO frequency correlates with $I_{\rm X}$, on timescales of days or longer the kHz QPO frequency can be the same even if $I_{\rm X}$ is a factor of $\sim 2$ different. This is the so-called parallel-track phenomenon (\citealt*{m1999}, \citealt*{v2001}). If the frequency of (one of) the kHz QPOs reflects the Keplerian frequency at the inner edge of the disc, and the position of the inner edge of the disc depends on $\dot M_i$ as in eq. \eqref{radius_mill}, the parallel-track phenomenon means that there cannot be a one-to-one relation between $\dot M_i$ and $I_{\rm X}$, and that the procedure that we described in the previous paragraph is incorrect. To overcome this problem, we proceeded as follows: Since the frequencies of the kHz QPOs appear to follow a random walk (\citealt{bmh2005}; \citealt{bv2012}), we generated values of $\dot M_i$ following a random walk\footnote{We restricted $\dot{M_i}$ to be between 1\% and 3\% $\dot{M}_E$ generating uniformly random distributed steps between $-10^{-4}$ and $10^{-4}$ $\dot{M}_E$.} we calculated $\ddot M_i$, and we computed $\dot \nu_{\rm sr}$ vs. $\nu_{\rm sr}$ using eq. (3) to (5). Following \citet*{mlp1998}, we took $h/r = 0.1$ and $v^r = 10^{-5} c$, we assumed a NS with a radius of 13 km and a mass of 1.9 $M_\odot$, and generated values of $\dot M_i$ between 0.01 and 0.03 $\dot M_E$. We further assumed that the Keplerian frequency at the sonic-point radius is the upper kHz QPO (\citealt*{mlp1998}; but see below), and converted the sonic-point Keplerian frequency to that of the lower kHz QPO using the relation $\nu_u = 0.673 \times \nu_l + 539$ Hz from \citet{bmh2005}, where $\nu_l$ and $\nu_u$ are, respectively, the frequencies of the lower and the upper kHz QPO. In Figure \ref{simultion_vs_data} we plot the observed relation $\dot \nu_{\rm QPO}$ vs. $\nu_{\rm QPO}$ (red points; these are the same data as in Figure \ref{speed64}) together with the relation between $\dot \nu_{\rm sr}$ and $\nu_{\rm sr}$ from this simulation (grey stars). From this Figure it is apparent that the simulation reproduces the observed data well. We note that although the range of $\dot M_i$ in our simulation is a factor of $\sim 3$ less than the range of luminosities spanned by 4U 1636--53 (e.g., \citealt{b3}), this may be accounted for if we adjust the values of $h/r$ and $v^r$ accordingly. Finally, we note that we can reproduce the data equally well if we assume that the frequency at the sonic-point radius is the lower kHz QPO if, in this case, the neutron star has a radius of 14 km and a mass of 1.7 $M_\odot$. In summary, assuming that the frequency of one of the kHz QPOs is equal to the Keplerian frequency at the inner edge of the accretion disc, the comparison between the observations and our simulations appears to lend support to the idea that the inner edge of the accretion disc is set by the interplay between viscosity and radiation drag. Our results, however, do not shed light about the mechanism that produces the kHz QPOs, or whether it is the lower or the upper kHz QPO the one that reflects the Keplerian frequency at the inner edge of the accretion disc. It remains to be seen whether other mechanisms (e.g., magnetic drag) can also reproduce the observations. \begin{figure} \begin{center} \includegraphics[width=80mm]{figure6.eps} \caption{Comparison between the absolute value of the frequency derivative of the lower kHz QPO frequency in 4U 1636--53 (red points) and simulated values of the frequency derivative of the lower kHz QPO using the \emph{sonic-point model} from \citet*{mlp1998} (grey stars). See section \ref{miller} for more details about the simulation.} \label{simultion_vs_data} \end{center} \end{figure} | 12 | 6 | 1206.0182 |
1206 | 1206.2839_arXiv.txt | We investigate the effect of a cosmological constant on the gravothermal catastrophe in the Newtonian limit. A negative cosmological constant acts as a thermodynamic `destabilizer'. The Antonov radius gets smaller and the instability occurs, not only for negative but also for positive energy values. A positive cosmological constant acts as a `stabilizer' of the system, which, in this case, exhibits a novel `reentrant behaviour'. In addition to the Antonov radius we find a second critical radius, where an `inverse Antonov transition' occurs; a series of local entropy maxima is restored. | Antonov's gravothermal instability is an important effect in gravitational thermodynamics \cite{Antonov,Bell-Wood,Padman,Chavanis}. It has served as the prime paradigm for an extensive research into the statistical mechanics of systems with long range interactions in different fields of physics \cite{Bell, Dauxois}, while phase transitions of self-gravitating systems is an active research field \cite{Chavanis2}. One may ask if and how this phenomenon depends on the vacuum background. A positive cosmological constant is nowadays one of the main candidates for dark energy \cite{Padman2,Harvey}, while negative cosmological constant has attracted a lot of attention due to the anti-de Sitter/Conformal Field Theory correspondence (AdS/CFT \cite{AdSCFT}). In addition, the stability properties of anti-de Sitter spacetime have become recently a subject of topical interest \cite{AdS1,AdS2}. We present here how gravothermal catastrophe is affected by a cosmological constant, i.e. by a de Sitter or anti-de Sitter vacuum. For convenience we call the non-relativistic limit of de Sitter \cite{Axenides} and anti-de Sitter spaces (usually called Newton-Hooke spaces \cite{NewtonHook}), just dS or AdS, respectively. \\ \indent The system under study in the original formulation \cite{Antonov,Bell-Wood} is a self-gravitating gas in the Newtonian limit bound by a spherical shell with insulating and perfectly reflecting walls with fixed energy (microcanonical ensemble) and fixed number of point particles (stars). An equilibrium state corresponding to an entropy maximum is called an `isothermal sphere'. Antonov proved that for such a system in the mean field approximation there is no global entropy maximum. However, there exist local entropy maxima (i.e. metastable states) for $E\cdot R > -0.335 GM^2$. This means that for positive energy $E$, there are isothermal spheres for every radius $R$, but for negative energy local entropy extrema exist only for radii smaller than a critical value $R_A$, we call Antonov radius. For $R > R_A$ and fixed negative energy there does not exist any equilibrium state. \\ \indent For a fixed negative energy and for radii smaller than the Antonov radius $R < R_A$, the equilibrium state may be stable (local entropy maxima) or unstable (saddle points) depending on the value of the ratio of core density $\rho_0 = \rho(0)$ versus the edge ratio $\rho_R=\rho(R)$. There is a critical value $(\rho_0/\rho_R)_{cr} = 709$ at which an instability sets in, i.e. for $\rho_0/\rho_R > 709$. A remarkable feature of gravity is that this instability in the microcanonical ensemble sets in when the specific heat goes from negative to positive values and not the other way around (non-equivalence of ensembles). \\ \indent We find that gravothermal catastrophe (described in the last two paragraphs) depends crucially on the vacuum background. AdS space destabilizes the system. Compared to the flat case, the instability sets in at smaller radius, at higher central density and occurs not only for negative but also for positive energies. In dS space the phenomenon is drastically altered. A series of equilibrium solutions is restored for large radii, indicating a type of `reentrant' behaviour. A metastable homogeneous solution that suffers a transition to Antonov instability is found, together with metastable states that do not suffer any transition to instability. In addition, centrally diluted and periodically condensed solutions are allowed. | 12 | 6 | 1206.2839 |
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1206 | 1206.5302_arXiv.txt | Cosmological surveys aim to use the evolution of the abundance of galaxy clusters to accurately constrain the cosmological model. In the context of $\Lambda$CDM, we show that it is possible to achieve the required percent level accuracy in the halo mass function with gravity-only cosmological simulations, and we provide simulation start and run parameter guidelines for doing so. Some previous works have had sufficient statistical precision, but lacked robust verification of absolute accuracy. Convergence tests of the mass function with, for example, simulation start redshift can exhibit false convergence of the mass function due to counteracting errors, potentially misleading one to infer overly optimistic estimations of simulation accuracy. Percent level accuracy is possible if initial condition particle mapping uses second order Lagrangian Perturbation Theory, and if the start epoch is between 10 and 50 expansion factors before the epoch of halo formation of interest. The mass function for halos with fewer than $\sim 1000$ particles is highly sensitive to simulation parameters and start redshift, implying a practical minimum mass resolution limit due to mass discreteness. The narrow range in converged start redshift suggests that it is not presently possible for a single simulation to capture accurately the cluster mass function while also starting early enough to model accurately the numbers of reionisation era galaxies, whose baryon feedback processes may affect later cluster properties. Ultimately, to fully exploit current and future cosmological surveys will require accurate modeling of baryon physics and observable properties, a formidable challenge for which accurate gravity-only simulations are just an initial step. | \let\thefootnote\relax\footnotetext{$^*$email: \texttt{ [email protected]}} In the vacuum energy dominated cold dark matter cosmological model \citep[hereafter $\Lambda$CDM,][]{Komatsuetal2011short}, large-scale structures form through the amplification of small density fluctuations via gravitational instability. At early times this amplification can be followed using linear perturbation theory of the general relativistic equations of motion for the field. At late times, owing to the nonlinearities in the equations, and after shell-crossing, the dynamics may only be accurately followed using numerical simulations. Overdense regions of the density field, whose dynamics have broken away from the evolution of the background space-time and have reached some state of virial equilibrium are commonly referred to as dark matter halos. The growth rate of large-scale structures is directly sensitive to the expansion rate of the Universe, and hence the cosmological parameters. One can show theoretically, through the excursion set formalism \citep{PressSchechter1974,Bondetal1991,ShethTormen1999}, that the number of halos is also sensitive to cosmological parameters, and importantly for future surveys, the presence of ``dark energy'' \citep{WangSteinhardt1998,Haimanetal2001,LimaHu2004,MarianBernstein2006,Cunhaetal2010,Courtinetal2011}. This forecast cosmological sensitivity has recently been verified through direct testing with $N$-body simulations \citep{SmithMarian2011}. The amount of cosmological information that can be extracted from cluster number counts is limited by our ability to detect signal-to-noise peaks in our observational survey -- \ie associate galaxies to groups, identify groups relative to an X-ray background noise level, etc. The lower that one can push the minimum detectable mass, the more cosmological information can be extracted from the survey. This comes under the proviso that one can accurately calibrate the true--observable mass relationship \citep{LimaHu2005,Marianetal2009,Rozoetal2009,Mandelbaumetal2010,OguriTakada2011,Anguloetal2012}. The numbers of rare halos are also sensitive to the level of non-Gaussianity in the primordial density field due to its effect upon the tail of extreme density fluctuations \citep{Matarreseetal2000,Marianetal2011a}. Cluster counts are also sensitive to the total neutrino mass \citep[\eg]{Wangetal2005,Carboneetal2012,Shimonetal2012} Thus, surveys that promise to accurately measure the evolution of the abundance of groups and clusters, also have the potential to help probe fundamental physics. Accurate theoretical predictions for the cluster mass function and its dependence on cosmology, are therefore essential to fully exploit next generation cluster surveys. Current cosmological constraints from clusters come from: \citet{Vikhlininetal2009,Vanderlindeetal2010short,Rozoetal2010, Sehgaletal2011short,Allenetal2011,Planck2011SZshort}. Over the next decade there will be a number of large surveys that will aim to strongly constrain the cosmological model through the abundance of clusters: in the X-ray there will be eROSITA \citep{Pillepichetal2012}, with the Sunyaev-Zel'Dovich method there will be Planck\footnote{www.rssd.esa.int/planck}, in the optical using the weak lensing method there will be DES\footnote{www.darkenergysurvey.org}, Euclid \citep{EUCLID2011} and LSST\footnote{www.lsst.org/lsst}. Several authors have estimated the requirements on the theoretical accuracy of the halo mass function to achieve the statistically limited constraints on cosmological parameters. \citet{Wuetal2010} point out that, in order to constrain time evolving dark energy models for DES, the theoretical mass function must be known with an accuracy $\lesssim0.5\%$. In this paper we address the question: What are the correct numerical parameters needed to achieve percent level accuracy in the mass function in a cosmological simulation? Large simulation volumes (whether by a single simulation or by multiple realizations) are able to reduce statistical uncertainties due to finite halo numbers. However, large absolute volumes are needed to reduce systematic and statistical errors associated with poor sampling of large-scale density modes \citep[e.g.][]{BarkanaLoeb2004,BaglaRay2005,PowerKnebe2006,Reedetal2007b,Crocceetal2010,SmithMarian2011}. Over the past decade, impressive statistical precision in the halo mass function has been achieved using suites of cosmological simulations \citep[][and others]{Jenkinsetal2001,Warrenetal2006,Reedetal2007b,Tinkeretal2008,Crocceetal2010,Ilievetal2012,Bhattacharyaetal2011,SmithMarian2011,Anguloetal2012,Alimietal2012,Watsonetal2012}. However, statistical precision does not imply accuracy, even when considering gravity-only simulations. Sources of systematic error include finite simulation volume, force resolution, mass resolution and discreteness effects, time-stepping, halo finding, initial condition particle mapping, and start redshift. Recent progress includes \citet{Crocceetal2010} and \citet{Bhattacharyaetal2011}, who each address many of the systematic uncertainties and determine a halo mass function with an estimated accuracy of $\sim2\%$ from a suite of large gravity-only cosmological boxes, though their results differ by significantly more than this for halos larger than $\sim 10^{15} \Msol$. Moreover, neither approach has taken into account the full covariance matrix of mass function estimates when deriving best fit parameters \citep{SmithMarian2011}. As a first step on the path toward producing an accurate mass function in the observational plane, we limit ourselves to demonstrating how percent level accuracy in gravity-only (\ie collisionless) simulations (wherein baryons are present but interact only via gravity) can be accomplished. We show how to set up initial conditions so that percent level accuracy can be achieved. We also isolate and test the run parameters that control force resolution (force softening and tree opening angle) and time-step size, allowing us to determine the required values to achieve percent level convergence. Finally, we ask: how many particles do we need to sample the halo mass distribution, in order to obtain a mass function accurate to better than $\lesssim1\%$. The paper breaks down as follows: in \S\ref{sec-techniques} we discuss setting up the initial conditions for the structure formation simulations and the parameters used to run the $N$-body codes. In \S\ref{sec-simulations} we describe the suite of $N$-body simulations performed and halo identification. In \S\ref{sec-mf} we present the results for the halo mass function and its convergence with simulation parameters. In \S\ref{sec-ic} we explore the variation of the halo mass function with the method for generating the initial conditions. We also make a comparison between the results obtained from two well known $N$-body codes. In \S\ref{sec-conother} we explore the convergence of the matter power spectrum and the 1-point probability density function of matter fluctuations. In \S\ref{sec-challenges} we discuss the remaining challenges of obtaining better than 1\% accurate mass functions from structure formation simulations. In \S\ref{sec-guidelines} we summarize our findings and draw up a set of guidelines for obtaining accurate gravity-only mass functions. | \label{sec-guidelines} In this paper we have explored the dependence of the mass function of dark matter halos on simulation run parameters and initial conditions. Our aim has been to perform convergence tests that will illuminate the path to obtaining percent level accuracy in this statistic. This will be a requirement for future cluster surveys of the Universe that aim to help constrain the nature of dark energy or dark gravity. In \S\ref{sec-techniques} we gave a brief overview of the simulation method, paying special attention to how one sets up initial conditions, either using Zel'Dovich approximation (1LPT) or 2LPT. We described the simulation codes that we have employed {\tt PKDGRAV} and {\tt Gadget-2}, with the former being the main code used throughout this study. In \S\ref{sec-simulations} we described the large suite of $N$-body simulations that we have performed to study these problems. All simulations were run at $N=1024^3$ and we covered two regimes: high redshift ($z=10$), small scale ($L=17.625\Mpc$) and low redshift ($z=0$) large scale ($L=2\Gpc$). In \S\ref{sec-mf} we explored the dependence of the mass functions on the simulation parameters and found the following: the resultant mass functions were rather insensitive to the choice of the tree-opening angle, provided $\Theta<0.7$; the results for halos resolved with fewer than $N\sim1000$ particles were sensitive to the choice of force softening, with larger values tending to increase the abundance of halos in this regime; results were fairly insensitive to the size of the adaptive time-step parameter and that 1\% converged results could be achieved for $\eta\lesssim0.15$. We also demonstrated that the use of anti-aliasing filters, such as the Hann filter, to set up initial conditions, can lead to $\sim30\%$ suppression in the abundance of halos resolved with $N\lesssim1000$ particles. We do not advocate the use of the Hann filter, since {\em there is no aliasing} in the initial conditions to correct. In \S\ref{sec-ic} we performed a detailed study of the impact of the choice of initial conditions on the mass function. We found that the results from simulations that are initialized with 1LPT converge very slowly as the start redshift is increased. The effect of too low a start redshift being the suppression of the formation of high mass halos. Furthermore, for the large box simulations, we also found simulations started at very high redshifts $z_{\rm i}\gtrsim200$ would fail to correctly follow the build up of structure due to the relative increase in numerical noise. Furthermore, 1LPT initial conditions exhibit ``false convergence'' with increasing start redshift. Simulations starting from 2LPT initial conditions proved to have very good convergence properties and for simulations that underwent 10-50 expansion factors, yielded percent level convergence in the halo mass function at the $1024^3$ resolution of our tests. We made a direct comparison of these results obtained from integration of the initial conditions with the tree-code {\tt ~PKDGRAV} with results from the Tree-PM code {\tt ~Gadget-2}, and found almost identical behaviour. However, a detailed comparison of the mass functions from the two codes revealed that {\tt Gadget-2} produced a $\sim10\%$ increase in the mass function for halos resolved with $N\lesssim10^2$ particles. These results extend and support the earlier findings of \citep{Crocceetal2006}. In \S\ref{sec-conother} we explored the convergence properties of two other statistics of the density field, namely the matter power spectrum and the 1-point probability density function (PDF) of matter fluctuations. We found that the simulation parameters that produced $\lesssim1\%$ convergence in the mass function would also lead to good convergence behaviour in these statistics. In addition, too high a start redshift for either 1LPT or 2LPT initial conditions would lead to systematic errors. On the other hand, the results from the simulations run with 2LPT initial conditions demonstrated excellent convergence behaviour. In summary, Table~\ref{table-recipe} presents a general recipe for the parameters needed for percent accuracy of the mass function within a gravity-only simulation using a tree code. Except for the tree opening-angle $\Theta$, which has some dependence on the specific tree used, all the other run parameters can be applicable to other tree codes. This list shows required values but is not complete. In future work, one would expect this table to be extended to include the following: if PM forces are used for large scale force computation, then parameters controlling their accuracy, such as the size of the PM grid should be included; multipole expansions are used to compute the tree forces, and different codes use different orders: which order is sufficiently accurate for our purposes? Also, there should be some entry associated with the parameters that control the halo finder (halo definition). Ultimately, inferring cosmological parameters from the cluster mass function will require a number of other issues to be solved relating to baryons and observable properties. Among the difficulties that baryons pose is the gravitational coupling of baryon processes to dark matter \citep{SomogyiSmith2010,vanDaalenetal2011}. Inferring observable properties from the simulations for comparison via mass-observable relations or by direct mock catalogs is a further complexity. Thus, percent level accuracy in numerical simulations represents a formidable challenge, but one that we must meet if future surveys of the Universe are to live up to their potential. | 12 | 6 | 1206.5302 |
1206 | 1206.4304_arXiv.txt | {We study the shapes of galaxy dark matter haloes by measuring the anisotropy of the weak gravitational lensing signal around galaxies in the second Red-sequence Cluster Survey (RCS2). We determine the average shear anisotropy within the virial radius for three lens samples: the `all' sample, which contains all galaxies with $19<m_{r'}<21.5$, and the `red' and `blue' samples, whose lensing signals are dominated by massive low-redshift early-type and late-type galaxies, respectively. To study the environmental dependence of the lensing signal, we separate each lens sample into an isolated and clustered part and analyse them separately. We address the impact of several complications on the halo ellipticity measurement, including PSF residual systematics in the shape catalogues, multiple deflections, and the clustering of lenses. We estimate that the impact of these is small for our lens selections. Furthermore, we measure the azimuthal dependence of the distribution of physically associated galaxies around the lens samples. We find that these satellites preferentially reside near the major axis of the lenses, and constrain the angle between the major axis of the lens and the average location of the satellites to $\langle \theta \rangle=43.7^{\circ}\pm0.3^{\circ}$ for the `all' lenses, $\langle \theta \rangle=41.7^{\circ}\pm0.5^{\circ}$ for the `red' lenses and $\langle \theta \rangle=42.0^{\circ}\pm1.4^{\circ}$ for the `blue' lenses. We do not detect a significant shear anisotropy for the average `red' and `blue' lenses, although for the most elliptical `red' and `blue' galaxies it is marginally positive and negative, respectively. For the `all' sample, we find that the anisotropy of the galaxy-mass cross-correlation function $\langle f-f_{45} \rangle=0.23\pm0.12$, providing weak support for the view that the average galaxy is embedded in, and preferentially aligned with, a triaxial dark matter halo. Assuming an elliptical Navarro-Frenk-White (NFW) profile, we find that the ratio of the dark matter halo ellipticity and the galaxy ellipticity $f_h=e_h/e_g=1.50_{-1.01}^{+1.03}$, which for a mean lens ellipticity of 0.25 corresponds to a projected halo ellipticity of $e_h=0.38_{-0.25}^{+0.26}$ if the halo and the lens are perfectly aligned. For isolated galaxies of the `all' sample, the average shear anisotropy increases to $\langle f-f_{45} \rangle=0.51_{-0.25}^{+0.26}$ and $f_h=4.73_{-2.05}^{+2.17}$, whilst for clustered galaxies the signal is consistent with zero. These constraints provide lower limits on the average dark matter halo ellipticity, as scatter in the relative position angle between the galaxies and the dark matter haloes is expected to reduce the shear anisotropy by a factor $\sim$2. } | \hspace{4mm} Over the last few decades a coherent cosmological paradigm has developed, $\Lambda$CDM, which provides a framework for the study of the formation and evolution of structure in the Universe. N-body simulations that are based on $\Lambda$CDM predict that (dark) matter haloes collapse such that their density profiles closely follow a Navarro-Frenk-White profile \citep[NFW;][]{Navarro96}, which is in excellent agreement with observations. Another fundamental prediction from simulations is that the haloes are triaxial \citep[e.g.][]{Dubinski91,Allgood06}, which appear elliptical in projection. This prediction of dark matter haloes, as well as many others concerning the evolution of their shapes \citep[e.g][]{Vera-Ciro11}, the effect of the central galaxy on the dark matter halo shape \citep[e.g.][]{Kazantzidis10,Abadi10,Machado10} and their dependence on environment \citep[e.g][]{Wang11}, remain largely untested observationally. \\ \indent Direct observational constraints on the halo ellipticities have proven to be difficult, mainly due to the lack of useful tracers of the gravitational potential. On small scales ($\sim$few kpc), halo ellipticity estimates have been obtained through the combination of strong lensing and stellar dynamics \citep[e.g.][]{VanDeVen10,Dutton11,Suyu11}, planetary nebulae \citep[e.g.][]{Napolitano10} and HI observations in late-type galaxies \citep[e.g.][]{Banerjee08,Obrien10}. On larger scales, the distribution of satellite galaxies around centrals has been used \citep[e.g.][]{Bailin08}, but such studies have only provided constraints for rich systems that may not be representative for the typical galaxy in the universe. \\ \indent Weak gravitational lensing does not depend on the presence of optical tracers and is capable of providing ellipticity estimates on a large range of scales (between a few kpc to a few Mpc). Therefore it is a powerful observational technique to study the ellipticity of dark matter haloes. In weak lensing the distortion of the images of faint background galaxies due to the dark matter potentials of intervening structures, the lenses, is measured. This has been used to determine halo masses \citep[e.g.][]{VanUitert11} as well as the extent of haloes. If galaxies preferentially align (or anti-align) with respect to the dark matter haloes in which they are embedded, the lensing signal becomes anisotropic. This signature can be used to constrain the ellipticity of dark matter haloes of galaxies \citep{Brainerd00,Natarajan00}. \\ \indent The core assumption in the weak-lensing-based halo ellipticity studies is that the orientation of galaxies and dark matter haloes are correlated; if they are not, the shear signal is isotropic and cannot be used to constrain the ellipticity of the haloes. The relative alignment between the baryons and the dark matter has been addressed in a large number of studies based on numerical simulations \citep[e.g.][]{VanDenBosch02,VanDenBosch03,Bailin05,Kang07,Bett10,Hahn10,Deason11}, in studies based on the distribution of satellite galaxies around centrals \citep[]{Wang08,Agustsson10} and in studies based on the ellipticity correlation function \citep{Faltenbacher09,Okumura09}. The general consensus is that although the galaxy and dark matter are aligned on average, the scatter in the differential position angle distribution is large. \citet{Bett11} examined a broad range of galaxy-halo alignment models by combining $N$-body simulations with semi-analytic galaxy formation models, and found that for most of the models under consideration, the stacked projected axis ratio becomes close to unity. Consequently, the ellipticity of dark matter haloes may be difficult to measure with weak lensing in practice. \\ \indent Knowledge of the relative alignment distribution is not only crucial for halo ellipticity studies, but also for studies of the intrinsic alignments of galaxies. Numerical simulations predict that the shapes of neighbouring dark matter haloes are correlated \citep[e.g.][]{Splinter97,Croft00,Heavens00,Lee08}. The shapes of galaxies that form inside these haloes may therefore be intrinsically aligned as well. Measuring this effect is interesting as it provides constraints on structure formation. Also, the lensing properties of the large-scale structure in the universe, known as cosmic shear, are affected by intrinsic alignments, and benefit from a careful characterization of the effect. Intrinsic alignments are studied observationally by correlating the ellipticities of galaxies as a function of separation; misalignments can significantly reduce these ellipticity correlation functions \citep[e.g.][]{Heymans04}. \\ \indent To date, only three observational weak lensing studies have detected the anisotropy of the lensing signal \citep{Hoekstra04,Mandelbaum_ell06,Parker07}. These studies have provided only tentative support for the existence of elliptical dark matter haloes, as they were limited by either their survey size and lack of colour information \citep{Hoekstra04,Parker07} or their depth \citep{Mandelbaum_ell06}. To improve on these constraints, we use the Red-sequence Cluster Survey 2 \citep[RCS2, ][]{Gilbank10}. Covering 900 square degree in the $g'r'z'$-bands, a limiting magnitude of $r'_{\rm{lim}}\sim 24.3$ and a median seeing of 0.7$''$, this survey is very well suited for lensing studies \citep[see][]{VanUitert11}. Using the colours we select massive luminous foreground galaxies at low redshifts. To investigate whether the formation histories and environment affect the average halo ellipticity of galaxies, the lenses are separated by galaxy type and environment, and the signals are studied separately. \\ \indent The structure of this paper is as follows. We describe the lensing analysis, including the data reduction of the RCS2 survey, the lens selection and the definition of the shear anisotropy estimators, in Section \ref{sec_analysis5}. We present measurements using a simple shear anisotropy estimator in Section \ref{sec_rat}, and use it to study the potential impact of PSF residual systematics in the shape catalogues. Various complications exist that might have altered the observed shear anisotropy, and in Section \ref{sec_multi} we study the impact of two of them: multiple deflection and the clustering of the lenses. The shear anisotropy measurements are shown and interpreted in Section \ref{sec_he}. We conclude in Section \ref{sec_concl5}. Throughout the paper we assume a WMAP7 cosmology \citep{Komatsu11} with $\sigma_8=0.8$, $\Omega_{\Lambda}=0.73$, $\Omega_M=0.27$, $\Omega_b=0.046$ and the dimensionless Hubble parameter $h=0.7$. The errors on the measured and derived quantities in this work generally show the 68\% confidence interval, unless explicitly stated otherwise. | } \hspace{4mm} We present measurements of the anisotropy of the weak lensing signal around galaxies using data from the Red-sequence Cluster Survey 2 (RCS2). We define three lens samples: the `all' sample contains all galaxies in the range $19<m_{r'}<21.5$, whereas the `red' and `blue' samples are dominated by massive low-redshift early-type and late-type galaxies, respectively. To study the environmental dependence of the lensing signal, we also subdivide each lens sample into an isolated and clustered part, and analyse them separately. \\ \indent We address the impact of several complications on the shear anisotropy measurements, including residual PSF systematics in the shape catalogues, multiple deflections, the clustering of lenses, and correlations between their intrinsic shapes. We run a set of idealised simulations to estimate the impact these might have on real data, and find them to be small, but not entirely negligible. We demonstrate that the impact of these complications can be reduced by a careful selection of the lens sample, i.e. low-redshift, massive and elliptical galaxies, as has been done in this work. \\ \indent We also measure the distribution of physically associated galaxies around the lens samples. We find that these satellites predominantly reside near the major axis of the lenses. The results of the `red' sample are in good agreement with previously reported values, whilst the constraints of the `all' and `blue' sample cannot be easily compared as they consist of a mixture of early-type and late-type galaxies. \\ \indent The shear anisotropy is quantified by the anisotropy of the galaxy-mass cross-correlation function, $\langle f-f_{45}\rangle$, and by the ratio of the projected dark matter halo ellipticity and the observed galaxy ellipticity, $f_h$. For the `all' sample we find that $\langle f-f_{45}\rangle=0.23\pm0.12$, and $f_h=1.50_{-1.01}^{+1.03}$ for an elliptical NFW profile, which for a mean lens ellipticity of 0.25 corresponds to a projected halo ellipticity of $e_h=0.38_{-0.25}^{+0.26}$ if the halo and the lens are perfectly aligned. Note that various studies indicate that this may not be the case. These constraints provide weak support that galaxies are embedded in, and preferentially aligned with, triaxial dark matter haloes. For isolated galaxies, the average shear anisotropy is larger than for clustered galaxies; for elliptical NFW profiles, we find $f_h=4.73_{-2.05}^{+2.17}$ and $f_h=0.90_{-1.15}^{+1.17}$, respectively. The decrease of the lensing anisotropy signal around clustered galaxies may be due to the stripping of dark matter haloes in dense environments. \\ \indent We do not detect a significant shear anisotropy for the average `red' lens. The shear anisotropy for the most elliptical `red' galaxies is marginally positive, but the pattern of the signal suggests that this is not the result of an alignment between the dark matter haloes and the galaxies. For the `blue' lenses, we find that the shear is marginally negative, with slightly increased significance for the most elliptical galaxies, suggesting an anti-alignment between the galaxy and the dark matter. Our measurements highlight the need for (photometric) redshifts in lensing studies. In order to reach sufficient signal-to-noise that enable competitive constraints on the shear anisotropy, we have to stack large numbers of galaxies that span a broad range in luminosities and redshifts. This smears out the shear anisotropy, and in the worst case the anisotropy might even average out. \\ \paragraph{Acknowledgements \\ \\} We would like to thank Elisabetta Semboloni for useful discussions and suggestions on the error estimation of the shear anisotropy measurements, Peter Schneider for valuable comments on the manuscript, and Rachel Mandelbaum for providing the table of the shear anisotropy of an elliptical NFW profile as well as for valuable comments on this work. HH and EvU acknowledge support from a Marie Curie International Reintegration Grant. HH is also supported by a VIDI grant from the Nederlandse Organisatie voor Wetenschappelijk Onderzoek (NWO). TS acknowledges support from NSF through grant AST-0444059-001, and the Smithsonian Astrophysics Observatory through grant GO0-11147A. MDG thanks the Research Corporation for support via a Cottrell Scholars Award. The RCS2 project is supported in part by grants to HKCY from the Canada Research Chairs program and the Natural Science and Engineering Research Council of Canada. \\ \indent This work is based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institute National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii. We used the facilities of the Canadian Astronomy Data Centre operated by the NRC with the support of the Canadian Space Agency. | 12 | 6 | 1206.4304 |
1206 | 1206.6137_arXiv.txt | For the first time, the kinematic evolution of a coronal wave over the entire solar surface is studied. Full Sun maps can be made by combining images from the \textit{Solar Terrestrial Relations Observatory} satellites, \textit{Ahead} and \textit{Behind}, and the \textit{Solar Dynamics Observatory}, thanks to the wide angular separation between them. We study the propagation of a coronal wave, also known as ``EIT'' wave, and its interaction with a coronal hole resulting in secondary waves and/or reflection and transmission. We explore the possibility of the wave obeying the law of reflection of waves. In a detailed example we find that a loop arcade at the coronal hole boundary cascades and oscillates as a result of the EUV wave passage and triggers a wave directed eastwards that appears to have reflected. We find that the speed of this wave decelerates to an asymptotic value, which is less than half of the primary EUV wave speed. Thanks to the full Sun coverage we are able to determine that part of the primary wave is transmitted through the coronal hole. This is the first observation of its kind. The kinematic measurements of the reflected and transmitted wave tracks are consistent with a fast-mode MHD wave interpretation. Eventually, all wave tracks decelerate and disappear at a distance. A possible scenario of the whole process is that the wave is initially driven by the expanding coronal mass ejection and subsequently decouples from the driver and then propagates at the local fast-mode speed. | Coronal waves are large scale bright fronts that propagate over the solar surface, with speeds ranging from 50 to over 700 km s$^{-1}$ \citep[e.g.][]{Thompson_09,Warmuth_11}. These features have traditionally been called ``EIT waves'', because they were first studied in detail \citep{Moses_97,Thompson_98,Thompson_99} with the Extreme Ultraviolet Imaging Telescope \citep[EIT;][]{Delaboudiniere_95}. It is more appropriate to call them ``EUV waves'' because they are observed in coronal extreme ultra-violet (EUV) wavelengths. ``EUV waves'' have been associated with similar phenomena in other wavelengths, such as chromospheric He I \citep{Vrsnak_02,Gilbert_04a,Gilbert_04b}, X-rays \citep{Khan_02,Narukage_02,Hudson_03,Warmuth_05,Vrsnak_06}, and radio wavelengths \citep{Klassen_00,Pohjolainen_01,Khan_02,Warmuth_04,Vrsnak_05,White_05,Vrsnak_06}. They are also associated with coronal mass ejections (CMEs), flares, and type II radio burst \citep{Biesecker_02,Cliver_05}. Recent reviews on this topic can be found in e.g. \citet{Wills-Davey_09} \citet{Gallagher_10}, \citet{Warmuth_10}, \citet{Zhukov_11}, and \citet{Patsourakos_12}. For the remainder of this paper we will refer to this phenomena as a ``coronal wave'' or ``wave front'', which is appropriate because of their wave nature. The interpretation of these waves has been a highly debated topic. Two broad scenarios have been proposed, either as a true wave (MHD wave propagating in the corona) or a pseudo-wave (a propagating front of an expanding structure, such as a CMEs outer envelope). Moreton waves \citep{Moreton_60} observed in H$\alpha$ have been interpreted as the chromospheric counterparts of fast-mode magneto-hydro-dynamic (MHD) coronal waves \citep{Uchida_68}. Hence, it was natural to associate Moreton waves and ``EIT waves'' \citep[e.g.][]{Warmuth_04}. Much evidence has been put forth for the MHD wave interpretation, as a fast-mode wave \citep[e.g.][]{Wills-Davey_99,Wang_00,Wu_01,Ofman_02,Patsourakos_09a,Patsourakos_09b,Schmidt_10}, a slow-mode wave \citep{Wang_09}, or soliton wave \citep{Wills-Davey_07}. The pseudo-wave interpretation proposes several mechanisms all associated with the large-scale structure of the associated CME. One such pseudo-wave is a propagating reconnection front \citep{Attrill_07a,Attrill_07b,vanDrielGesztelyi_08}, in which as the CME is expanding it interacts with the smaller scale loops in the vicinity via reconnection to produce heat and enhanced EUV intensity. Another possible pseudo-wave mechanism is the disk projection of large-scale current shells that surround the expanding flux-rope core of a CME \citep{Delannee_08}. Joule heating and/or plasma compression can occur at these shells and lead to enhanced intensities. Yet another pseudo-wave mechanism is the plasma compression due to the successive openings of the overlying magnetic fields as the CME is expanding \citep{Chen_02,Chen_05}. \citet{Patsourakos_09b} compiled a list of several observational tests for comparing the wave and pseudo-wave explanations, upon applying these test to one event they concluded in favor of a fast-mode wave interpretation and state that this wave is ``probably'' triggered by the expansion of the loops associated with the CME, in other words that it is driven by the CME. Finally, another point of view takes on a hybrid of wave and pseudo-wave interpretations \citep{Chen_02,Chen_05,Zhukov_04,Cohen_09,Chen_11,Downs_11,Cheng_12}. In this hybrid view a CMEs outer envelope is the pseudo-wave that can drive a fast-mode wave ahead, which subsequently evolves freely over the solar surface once the CME has propagated sufficiently away from the Sun. This emerging hybrid view shows the most promise for a complete description of coronal waves as it can account for most (if not all) observational features. The controversy seems to stem largely from ambiguities in the observed features due to limited spatial or temporal sampling. Recent ultra high cadence observations show clearly the formation of a shock in the EUV corona \citep{Liu_10, Kozarev_11,Ma_11,Cheng_12} providing strong support for the fast-mode interpretation. Further, the reflection of coronal waves at coronal hole (CH) boundaries is especially difficult to reconcile with a pseudo-wave mechanism. The first reported observation of a reflection of a coronal wave was associated with the event of 2007 May 19 \citep{Long_08,Veronig_08,Gopalswamy_09}, and more recently the 2011 June 7 event \citep{LiT_12}. \citet{Gopalswamy_09} studied the 2007 May 19 coronal wave event in detail and reported the reflection of the primary wave in three directions from a nearby CH west of the flaring active region, and possibly a fourth by a coronal hole to the south. To perform this analysis they used running differenced images to show this effect (the same technique was also used by \citet{LiT_12} to show an example of reflection). In the technique of running difference each image in a sequence of images is subtracted from the previous image to enhance moving features or intensity variations. The use of this technique caused some controversy because it can produce misleading wave-like effects. Upon reanalysis of this event, \citet{Attrill_10} showed that using running difference caused such misleading effects, or ``optical illusions'', and further comparison with a base differencing technique concluded that the ``wave'' does not in fact reflect from the coronal hole boundary. In base difference one selects a pre-event image and systematically subtracts that image from all the images in the sequence. \citet{Zhukov_11} comments that using this technique is also flawed in the case of a wave reflected back towards the eruption site because the background intensity over which the reflected wave propagates may have changed due to EUV dimming making the wave more difficult to detect. To explain what \citet{Gopalswamy_09} reports as ``reflection'' \citet{Attrill_10} invokes hot plasma channelled along coronal loops in the opposite direction of the ``wave'', a secondary event (presumably occurring very soon after the initial event, as suggested by the associated double-CME event), and/or a rotation of the ``wave'' due to local inhomogeneities of the fast-mode MHD wave speed. Therefore, the question of reflection remains open, especially since the techniques used to identify reflection have been called to question. Here we provide new observations of an event in which reflection could be a possible explanation, as the trajectory of the reflected wave seems to obey the law of reflection. In one example we show that there is a spatial and temporal correlation between the wave, which appears to have reflected from the CH boundary, and the oscillation of a loop arcade at the CH boundary aligned with the trajectory of the reflected wave. The loops in the arcade are triggered to oscillate one by one as the primary (or incident) wave reaches each individual loop thereby giving an appearance that the oscillation of the first loop triggers the oscillation of the next loop and so forth in what appears to resemble a cascade. However, the observations cannot conclusively determine whether the wave is in fact a true reflection or launched by the loop arcade. Excitation of coronal loop oscillations by the primary coronal wave has been reported in the literature \citep[e.g.][]{Wills-Davey_99,Ofman_02,Hudson_04,Ofman_07,Aschwanden_11}, though what is novel about this observation is the fact that an entire loop arcade oscillates and that the cascading speed was found to be related to the oblique angle at which the primary wave reached the arcade. In addition, we report the first detection of wave transmission through a CH, as previous observations seemed to indicate that coronal waves were stopped at the boundaries of CHs\citep[e.g.][]{Thompson_98,Veronig_06}. Both reflection and transmission have been predicted in recent MHD simulations of fast-mode coronal waves \citep{Schmidt_10}. \citet{Schmidt_10} identify in their simulations ``secondary waves'', which seem to describe both reflection and transmission waves. For the sake of generality, we use the same approach and define a ``secondary wave'' as any other observed wave induced (either directly or indirectly) by the primary wave including apparent reflected and transmitted waves. For example, we call the wave triggered by the cascading loop arcade at the CH boundary a secondary wave. Our paper is organized as follows. In \S~\ref{sec_obs}, we introduce the multi-view point and multi-wavelength space based EUV observations of the Sun and the full Sun EUV maps we use for analyzing the wave propagation on a global scale. In \S~\ref{sec_results}, we discuss our kinematic analysis of the wave reflection and transmission. We conclude in \S~\ref{sec_conclusion}. | We have demonstrated the wave nature of the coronal wave observed during the 2011 February 15 event by explicitly identifying and analyzing observations of secondary waves, which can be interpreted as reflection from and/or transmission through a CH. The asymptotic velocities, interpreted as reaching the local fast-mode speed, for all the wave ground track slices presented are comparable with each other (see table \ref{table_velo}). The difference in the asymptotic speeds of the ``reflected'' waves, slice ``A'', (380 vs. 280 km s$^{-1}$), of the meridional track toward the north, slice ``B'', (260 km s$^{-1}$) and of the SW track (280 km s$^{-1}$), that crossed the CH, seem to reflect the variations of the local fast-mode speed in the corona and are in agreement with global MHD models \citep[e.g.][]{Schmidt_10,Zhao_11}. Furthermore, these differences may provide good diagnostics for coronal seismology \citep[e.g.][]{Yang_10,West_11}, but this topic is outside of the scope of this paper. Our results are also consistent with other authors, for example \citet{Warmuth_11} showed, based on a statistical study, that coronal waves with initial fast velocities ($v \ge 320$ km s$^{-1}$) show the greatest deceleration and attain final velocities between 200-300 km s$^{-1}$, \citet{Veronig_10} found an asymptotic velocity of a coronal wave far from its origin of $\sim280$ km s$^{-1}$, and \citet{Long_08,Long_11} found deceleration in the events that they studied. This can be explained if the coronal wave were in fact a freely-propagating fast-mode MHD wave \citep[e.g.][]{Warmuth_04,Veronig_10,Long_11,Warmuth_11}. \citet{Warmuth_11} have proposed that for fast events ($v \ge 320$ km s$^{-1}$) the physical nature of the coronal wave can be explained as being initially a large-amplitude nonlinear wave and/or shock, presumably driven by the CME, that subsequently evolves (and decelerates) to a linear fast-mode wave propagating at the characteristic wave speed. This interpretation is agreement with the hybrid view of coronal waves, in which the pseudo-wave can explain the outer envelope of the CME that drives a fast-mode wave that is left freely propagating close to the solar surface once the CME has propagated radially away. It is this fast-mode wave that we are referring to in this paper. We have reported on features which appear to obey the law of reflection. Close inspection of the reflection to the east into the FOV of EUVI-B shows that the secondary wave observed may have been triggered by a cascading loop arcade. To the west, into the FOV of EUVI-A, a secondary wave propagates along a path that seems to obey the law of reflection. But, is what is seen to the west in fact an observation of reflection? Or a secondary wave launched by the active region or the CH resonating? It is difficult to tell with these observations because that region there are bright active region loops and the cadence of EUVI-A is not sufficient to resolve fine temporal structure in that region. Our work presented a case of reflection in which a different approach was taken by evoking the law of reflection and ray tracing possible trajectories of the wave towards and reflecting from the CH boundary. We used both running and base ratio stack plots as appropriate. While the reflection is a strong argument for the wave nature of coronal waves, it has been challenged as artifact of the data analysis method (\citet{Attrill_10} but see counter arguments in \citet{Zhukov_11}). When it comes to these techniques the primary problem arrises when analyzing a stack plot that is placed between the active region that launched the coronal wave and reflecting CH boundary surface. Two such possible effects may cause misleading effects to be seen in the stack plots. The first could be hot plasma, density enhancement, and/or a transverse/kink mode wave traveling along the loops aligned between the active region that launched the coronal wave and the CH boundary \citet{Attrill_10,Attrill_09}. The second could be loops triggered to oscillate along the CH boundary resulting in a signature in the stack plot that appears like reflection. Of course we must not be to hasty and dismiss the possibility that the wave did in fact reflect, the point we are trying to make is that each case of reflection should be considered carefully. The transmitted wave toward the SW can be identified in both running and base ratio images without ambiguity because it propagates across ``pristine'' quiet Sun. This is the first time, to our knowledge, that such a feature is identified in observations. The observation of the coronal hole crossing (or transmission) is a very strong argument for the wave nature of these EUV disturbances. Until now, it has been reported based on empirical evidence that coronal waves do not traverse coronal holes \citep[e.g.][]{Thompson_99,Veronig_06,Attrill_07b,Ma_09}. However, transmission has been discussed in the simulations of \citet{Schmidt_10}, who also discuss it in terms of waves launched by the ``resonating'' of the CH. Nevertheless we think that the same effect (e.g. transmission) is occurring in this event (either by resonance is not clear nor could it be determined based on these observations alone). Furthermore, we do see secondary waves launched all along the boundary of that CH (see provided movies), even beyond the reach of the original coronal wave, meaning that those waves would have had to cross the CH. But why it hasn't been observed before? We think that the size of the coronal hole is likely the determinant factor. If the coronal hole is large then the resonance may die out before it reaches the side opposite to the wave impact. On the other hand, if a hole is small the resonance may be sufficient to launch the secondary wave on the opposite side or all around the hole as suggested by \citet{Schmidt_10} and seen in this event as well as (partially) seen in the 2007 May 19 event \citep{Gopalswamy_09}. The reflection and transmission observations are hard to reconcile with a pure pseudo-wave interpretation and our results provide further evidence in favor for the wave nature of coronal waves. But note that it is the fast-mode wave that we are referring to originally driven by the CME, whose outer envelope may be described by the pseudo-wave. It may be the case that for weaker events a fast-mode wave is not triggered or is triggered but has an amplitude that does not produced enough heating to be detected, in these cases a pure pseudo-wave interpretation may be appropriate as the CME and its outer envelope may be the only signatures observed. In other similar events to the one studied here: the 2010 July 27 event reported by \citet{Chen_11}, the 2011 June 7 event reported by \citet{Cheng_12}, and the 2010 June 13 event reported by \citet{Downs_12}, it is shown that coronal wave events are a composite phenomenon comprised of a CME that drives a fast-mode wave, giving evidence to the hybrid interpretation. Based on those results what we have presented here is a detailed account of the wave component, initially driven by the CME, and that a pseudo-wave component should exist but the point at which one can distinguish between the two is not readily possible only to say that it should be early on in order to account for the secondary wave effects presented. We estimate that for this event the separation should have occurred some time before 02:00 UT based on the fact that the transmission is observed to occur after that time, and transmission is purely a wave effect. A projection effect of the CME material over the CH is the most obvious suggestion for the EUV emission inside the CH. Which could be the projection of deflected coronal material by the CME towards the south. But the CME material propagates away from the CH (towards the ecliptic plane) and its southern flank stops at the boundary \citep{Schrijver_11}. Even if we assume that the expanding material is somehow tilted to the SW (as seen from Earth), it is difficult to account for the appearance of a propagating low coronal feature outside of the CH to the SW (seen in the FOV of EUVI-A). The CME field lines cannot penetrate through the CH. No such observations has even been reported in countless limb observations of CMEs. The reflection is an additional problem. There is no secondary CME, we use both running and base ratio images and we do not track the reflection back through the incident path (as was done by \citet{Gopalswamy_09} and \citet{LiT_12}), so there is no chance for confusion. Secondary waves can also be seen in AIA observations of the CME and coronal wave event on 2011 June 7 \citep{LiT_12} and a transmitted wave can be seen in the 2011 February 24 CME. Instead, the fast-mode interpretation neatly accounts for all observations. Wave reflection is expected due to the sharp change in plasma parameters at the coronal hole/quiet sun interface as has been simulated \citep[][]{Schmidt_10} and reported observationally \citep{Long_08,Gopalswamy_09,LiT_12}, and wave transmission is predicted by MHD simulations \citep[][]{Schmidt_10}. Moreover oscillations of loops have also been reported as due to primary coronal wave fronts \citep[e.g.][]{Aschwanden_11}, which, as was shown here, can trigger secondary waves interpreted as being fast-mode MHD. In summary, we have undertaken a detailed analysis of the global kinematics of an EUV wave using the full Sun coverage afforded by the STEREO and SDO missions and have found several important results: \begin{itemize} \item Secondary wave effects attributed to the interaction of the primary coronal wave with the CH were described, including reflection and the first detection of wave transmission though a CH. \item We report detailed observations of the interaction of a coronal wave with a loop arcade that straddles the boundary of a CH, which subsequently launches a secondary wave, whose trajectory resembles the law of reflection. Such observations have not been reported before and suggest an origin to secondary waves that appear as reflections launched by CHs. \item Our observations and kinematic analysis is fully consistent with MHD simulations of fast-mode waves and suggests that these waves are initially driven, presumably by the expanding CME, but eventually relax to a freely propagating wave traveling at the local fast-mode speed. \item Our analysis shows that the quite Sun MHD wave speed can have modest variations. \item We also report the first global measurements of the angular extent of EUV waves, which reached at least or approximately $200^\circ$ before becoming too diffuse to be detected. \end{itemize} | 12 | 6 | 1206.6137 |
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